VLA Observational Status Summary 2025A

VLA capabilities October 2024 - May 2025

Introduction

Purpose of Document, Older Versions of the OSS

This Observational Status Summary (OSS) summarizes the instrumental status of the Karl G. Jansky Very Large Array (VLA) for the A-configuration, for the observing period 21 February 2025 through 25 August 2025 (Semester 2025A), and should be used when preparing proposals for the 31 January 2024 deadline. Multi-configuration proposals that include this configuration may also be submitted. Additionally, proposals requesting only configurations that will fall in future semesters may be submitted if the Principal Investigator is a graduate student. NRAO offers this service to provide scientific and technical feedback for students, and to provide them with an opportunity to re-submit their proposals for their principal semester with this information in hand.

The OSS is intended as a ready reference for those contemplating use of the VLA for their astronomical research. The information is in summary form; those requiring greater detail should use the NRAO Helpdesk, or refer to the manuals and documentation listed in Documentation. Most of the information contained here, and much more, is available through the VLA science web page and the companion VLBA science web page. For capabilities offered in previous semesters, we refer to our overview of all OSS versions available online.

The VLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for its efficient use. Although the NRAO strives to make using the VLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first–time users are encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. For more details, refer to the Visiting the DSOC and VLA page. The VLA is an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it.

 

Overview of the VLA

The Karl G. Jansky Very Large Array (VLA) is a 27–element interferometric array, arranged along the arms of an upside-down Y, which produces images of the radio sky at a wide range of frequencies and resolutions. The VLA is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico.

The basic data produced by the VLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array's elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing, commonly called phased array, allows operation of the array as a single element through coherent summation of the individual antenna signals. This mode is most commonly used for Very Long Baseline Interferometry (VLBI) observing and for observations of rapidly varying objects, such as pulsars.

The VLA can vary its resolution over a range exceeding a factor of ∼50 through movement of its component antennas. There are four basic antenna arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A, respectively. For details about antenna positions in the various configurations we refer to the stations position file (pdf version).

The VLA completes one cycle through all four configurations in an approximately 16 month period. Consult the Configuration Plans and Proposal Deadlines page or recent NRAO and AAS newsletters for current and up-to-date configuration schedules and associated proposal deadlines. Refer to the Guide to Proposing for the VLA for information on how to submit an observing proposal.

Observing projects on the VLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours with only one or, perhaps, a few target sources. However, since the VLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See the section on Snapshots for more detail.

All VLA antennas are outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers cover the frequency ranges of 1–2 GHz (L-band), 2–4 GHz (S-band), 4–8 GHz (C-band), 8–12 GHz (X-band), 12–18 GHz (Ku-band), 18–26.5 GHz (K-band), 26.5–40 GHz (Ka-band), and 40–50 GHz (Q-band). Additionally, all antennas of the VLA have receivers for lower frequencies, enabling observations at P-band (200–500 MHz). These low frequency receivers also work at 4-band (54–86 MHz), and new feeds have been deployed on all VLA antennas to observe at this frequency range.

The VLA correlator is both powerful and flexible. Details of the correlator configurations being offered for VLA science are described in the WIDAR Section of this document. It is important to realize that the VLA correlator is fundamentally a spectral line correlator and that even continuum observations are done in a wide-band mode with many channels.

Offered VLA Capabilities during the Next Semester

The Call for Proposals

The most recent Call for Proposals summarizes the General Observing (GO) capabilities being offered for the Karl G. Jansky Very Large Array (VLA).

In addition to these general capabilities, NRAO continues to offer shared risk observing options for those who would like to push the capabilities of the VLA beyond those offered for general use. These are the Shared Risk Observing (SRO) and Resident Shared Risk Observing (RSRO) programs.

Details about what is being offered for each program are given below. If you have any questions or problems with any link or tool, please submit a ticket through the NRAO Helpdesk.

Considering the lack of hybrid configurations after semester 2016A, guidelines on how to substitute such configurations with the use of principal array configurations are presented in the Array Configurations section of the Guide to Proposing for the VLA.

 

General Observing (GO) and Shared-Risk Observing (SRO)

Summary of Capabilities

As described in the Call for Proposals, the VLA offers continuous frequency coverage from 1–50 GHz in the following observing bands: 1–2 GHz (L-band); 2–4 GHz (S-band); 4–8 GHz (C-band); 8–12 GHz (X-band); 12–18 GHz (Ku-band); 18–26.5 GHz (K-band); 26.5–40 GHz (Ka-band); and 40–50 GHz (Q-band). Both single pointing and mosaics with discrete, multiple field centers will be supported under General Observing (GO). In addition to these, all VLA antennas are equipped with 224–480 MHz (P-band) and 54–86 MHz (4-band) receivers near the prime focus. Data rates of up to 60 MB/s (216 GB/hour) will be available to all users as GO, combined with correlator integration time limits per band and per configuration, as described in the Time Resolution and Data Rates section. Limitations on frequency settings and tuning ranges are described in the Frequency Bands and Tunability section.

The GO capabilities being offered are:

Capability Description
8-bit samplers
  • Standard full polarization default setups for:
    • 2 GHz bandwidth continuum observations at S/C/X/Ku/K/Ka/Q bands (16 × 128 MHz subbands)
    • 1 GHz bandwidth continuum observations at L-band (16 × 64 MHz subbands)
    • 256 MHz bandwidth continuum observations at P-band (16 × 16 MHz subbands)
    • 12 MHz bandwidth Stokes I continuum observations only* at 4-band (3 x 4 MHz subbands)
    • Dual 4/P-band for Stokes I continuum observations only*
  • Flexible setups for spectroscopy using two independently tunable, 1 GHz baseband pairs, each of which can be split into up to 32 flexibly tunable subbands
  • Single, dual, and full polarization products for non-default setups

*Note: 4-band and dual 4/P-band observations are offered for Stokes I continuum only using standard full polarization default setups. Polarization, spectral-line, or the use of non-standard setups, should be submitted as a RSRO proposal.
3-bit samplers
  • Standard full polarization default setups for:
    • 8 GHz bandwidth continuum observations at K/Ka/Q-bands
    • 6 GHz bandwidth at Ku-band
    • 4 GHz bandwidth at C/X-bands
  • Flexible setups for spectroscopy using four independently tunable, 2 GHz baseband pairs, each of which can be split into up to 16 flexibly tunable subbands
  • Single, dual, and full polarization products for non-default setups
Mixed 3-bit and 8-bit samplers
  • Allows more flexibility for simultaneous continuum and high-resolution spectral line observing

Subarrays

  • Up to 3 independent subbarrays using standard 3-bit continuum setups, or a mix of standard 3-bit and standard 8-bit continuum setups, and up to 3 independent subarrays with changing standard continuum setups in a given subarray (e.g., to perform reference pointing at X-band for high frequency observations).

Y27 or Y1 for VLBI

Solar observing

  • All solar observing except the L-band reverse-coupled system*

On-The-Fly Mosaicking (OTF)

  • P-, L-, S-, and C-bands only, using linear interpolations in Equatorial Coordinates; no subarrays

Pulsar

  • Phase-binned and coherent-dedispersion (YUPPI) pulsar observing, except 4-band YUPPI and combining YUPPI mode with VLBI recording, which are Resident Shared Risk Observing (RSRO) capabilities.

*Note: The VLA L-band (1-2 GHz) has a special signal path (the "reverse coupler" path) that allows coherent radio bursts to be observed without saturating the system, as the brightest of these solar bursts can exceed 105 solar flux units, or 109 Jy. This signal path has not yet been fully commissioned and is therefore not yet available under GO.

SRO capabilities can be set up via the Observing Preparation Tool (OPT) and run through the dynamic scheduler without intervention, but are not as well tested as GO capabilities. Data rates higher than 60 MB/s (216 GB/hour) and up to 100 MB/s (360 GB/hour) are considered SRO. A summary of the SRO capabilities being offered are:

  • On-the-Fly (OTF) mosaicking for X-, Ku-, K-, Ka-, and Q-bands (used when each pointing on the sky is on the order of several seconds or less), but not using subarrays.
  • Wideband VLA for VLBI: Enables recording of VLA WIDAR continuum-mode correlations during VLA phased array (Y27) VLBI observations. Currently, this only supports standard VLA 8-bit continuum modes with a 2-GHz bandwidth. See the VLBA Call for Proposals for more details.
  • eLWA: Joint LWA and VLA 4-band observations using a single 8 MHz subband centered at 76 MHz, and 4-bit VDIF output. Note: During semester 2025A, the LWA is expected to be undergoing infrastructure upgrades and availability of the telescopes (LWA1 and LWA-SV) may be limited.  Those interested in using this mode should contact Greg Taylor at gbtaylor@unm.edu for more details.

We expect that most SRO programs will have no or only minor problems that can be corrected quickly. If an SRO program fails, however, and it becomes clear that detailed testing with additional expertise is needed, then the project must make an experienced member from their team available to help troubleshoot the problem. In some cases, this may require the presence of that experienced member in Socorro. If adequate support from the project is not given, then the time on the telescope will be forfeited. The additional effort is to be determined based on discussions with the NRAO staff and management and the project team.

The guidelines for General and Shared Risk observing proposals, along with information about tools and other advice, can be found in the VLA Proposal Submission Guidelines.

Resident Shared Risk Observing (RSRO)

Summary of Capabilities

The VLBA Resident Shared Risk Observing (RSRO) program provides users with early access to new capabilities in exchange for a period of residency in Socorro to help commission those capabilities.

RSRO proposals should be submitted using the NRAO Proposal Submission Tool in response to a regular proposal call. The proposal should include a scientific justification, as for normal proposals, which will be peer reviewed as part of NRAO's time allocation process. Selecting "VLA RSRO" from the "Observing Mode" menu on the Resources page makes an "RSRO Comments" text-entry facility available for describing the technical resources required. A description of the personnel who will be involved in the effort along with their expertise and availability should also be included in the technical justification.

We emphasize the "shared risk" nature of the RSRO program. Since observers will be attempting to use capabilities under development and in the process of being commissioned, NRAO can make no guarantee of the success of any observations made under this program, and no additional commitment is made beyond granting the hours actually assigned by the peer review process.

Proposals for any area of user interest bit offered under GO or SRO are welcome. Here, we provide some examples of capabilities that are being utilized in recent RSRO proposals.

  • Correlator dump times shorter than 50 msec, including integration times as short as 5 msec for transient detection, or data rates above 100 MB/s. In order to reduce the data rate, frequency averaging in the correlator may be utilized in RSRO proposals;
  • YUPPI pulsar mode combined with VLBI recording;
  • Subarray observations with setups other than the default continuum setups, or observations with more than 3 subarrays;
  • Currently, OTF observing is implemented as linear interpolations in Equatorial Coordinates (i.e., RA/Dec). This can be expanded to allow using stripes linear in Galactic coordinates (l,b), as well as more complex patterns other than linear in RA and Dec, such as Rosetta or Spiral patterns. Note that these must still adhere to the restrictions of the OTF mode under General Observing, i.e., using the full array below 8 GHz (up to C-band), and no subarrays.

The guidelines for Resident Shared Risk Observing proposing, along with requirements and considerations, can be found in the VLA Proposal Submission Guidelines.

 

Commensal Observing Systems at the VLA

There are three commensal systems on the VLA that may take data at the same time as your proposed observation. The first is the VLITE system, which will take data at P-band during regular observations that use bands other than P-band. Hence, VLITE is turned off by default during P-band or dual 4/P-band observations. The VLITE system is deployed on up to eighteen VLA antennas. Observers wishing to gain access to the commensal VLITE data taken during their VLA observations should follow the instructions on the VLITE web page for doing so. The second is the realfast system, which takes data at very fast dump rates in an effort to detect Fast Radio Bursts (FRBs). This system is fully commissioned for observing at L- through X-bands, in parallel with standard continuum correlator configurations. The third commensal system, COSMIC SETI, enables the search for extraterrestrial intelligence (SETI) using the VLA, and collects data during unconflicted PI science observations. For information about commensal observing see the Commensal Observing with NRAO Telescopes page.

To report errors or problems encountered in any link or while using any NRAO tool listed here, please submit a ticket through the NRAO Helpdesk.

Performance of the VLA during the Next Semester

Resolution

Resolution

The VLA's resolution is generally diffraction-limited, and thus is set by the array configuration and the observing frequency. Like all synthesis arrays, the VLA is sensitive only to structures on a range of angular scales between the diffraction limit (the smallest angular scale detectable) and a "Largest Angular Scale" (which depends on the fringe spacing formed by the shortest baselines in the configuration). For emission structures smaller than the diffraction limit (θ ∼ λ/Bmax), the VLA acts like a single-dish instrument—the resulting image is smoothed to the resolution of the array. For emission structures larger than the detectable range, the VLA is simply blind to the emission; this is a limitation unique to interferometers. No subsequent processing can fully recover the missing information from these large scales. It can only be obtained by observing in a more compact VLA array configuration or with data from an instrument that is sensitive to the missing angular scales, such as a large single dish or a compact array of smaller antennas.

Table 3.1.1 displays the VLA's resolution and the scale at which severe attenuation of large-scale structure occurs. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power; the resolution element) for the central frequency for each band, and the largest angular scale of detectable emission.

Table 3.1.1: Configuration Properties
ConfigurationABCD
Bmax (km1) 36.4 11.1 3.4 1.03
Bmin (km1) 0.68 0.21 0.0355 0.035
Band Synthesized Beamwidth θHPBW(arcsec)1,2,3
74 MHz (4) 24 80 260 850
350 MHz (P) 5.6 18.5 60 200
1.5 GHz (L) 1.3 4.3 14 46
3.0 GHz (S) 0.65 2.1 7.0 23
6.0 GHz (C) 0.33 1.0 3.5 12
10 GHz (X) 0.20 0.60 2.1 7.2
15 GHz (Ku) 0.13 0.42 1.4 4.6
22 GHz (K) 0.089 0.28 0.95 3.1
33 GHz (Ka) 0.059 0.19 0.63 2.1
45 GHz (Q) 0.043 0.14 0.47 1.5
Band Largest Angular Scale θLAS(arcsec)1,4
74 MHz (4) 800 2200 20000 20000
350 MHz (P) 155 515 4150 4150
1.5 GHz (L) 36 120 970 970
3.0 GHz (S) 18 58 490 490
6.0 GHz (C) 8.9 29 240 240
10 GHz (X) 5.3 17 145 145
15 GHz (Ku) 3.6 12 97 97
22 GHz (K) 2.4 7.9 66 66
33 GHz (Ka) 1.6 5.3 44 44
45 GHz (Q) 1.2 3.9 32 32
These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith.
Notes:
1. Bmax is the maximum antenna separation, Bmin is the minimum antenna separation, θHPBW is the synthesized beam width (FWHM), and θLAS is the largest angular scale structure visible to the array.
2. The listed resolutions are appropriate for sources with declinations between −15 and +75 degrees.
3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θHPBW. The values for snapshots are about 1.3 times the listed values.
4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations, the quoted numbers should be divided by two.
5. For the C configuration, an antenna from the middle of the north arm is moved to the central pad N1. This results in improved imaging for extended objects, but may slightly degrade snapshot performance. Note that although the minimum spacing is the same as in D configuration, the surface brightness sensitivity and image fidelity to extended structure is considerably inferior to that of the D configuration.

The following figure is a graphical representation of the synthesized beamwidths for natural and robust weighting for the four main array configurations between 1 and 50 GHz. Also available are synthesized beamwidth figures for the low frequency (1–12 GHz) and the high frequency (12–50 GHz) receiver bands.

Sensitivity

Sensitivity

The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by:

where:

- SEFD is the system equivalent flux density (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEFD indicate more sensitive performance. For the VLA's 25–meter paraboloids, the SEFD is given by the equation SEFD = 5.62TsysA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band.
- ηc is the correlator efficiency (~0.93 with the use of the 8-bit samplers).
- npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in RCP or LCP.
- N is the number of antennas.
- tint is the total on-source integration time in seconds.
- Δν is the bandwidth in Hz.

Figure 3.2.1 shows the SEFDs as a function of frequency used in the VLA exposure calculator for those Cassegrain bands currently installed on VLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Figure 3.2.2 shows the SEFDs as a function of frequency for the P-band; these measurements are based on imaging of a field far from the galactic plane. Table 3.2.1 gives the SEFD at some fiducial VLA frequencies.

 Figure 3.2.1: SEFD used in the Exposure Calculator for the VLA. Left: The system equivalent flux density as a function of frequency for the L, S, C and X-band receivers. Right: The system equivalent flux density as a function of frequency for the Ku, K, Ka, and Q-band receivers. SEFDs at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere and were determined under good conditions. 

Figure 3.2.2: The SEFD used in the VLA Exposure Calculator as a function of frequency for the P-band receiver

Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical:

  • For the more commonly used robust weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values.
  • Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 3.2.1 shows the confusion noise in D configuration assuming robust weighting (see Condon et al. 2012, ApJ, 758),  which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 3.2.1; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This confusion primarily affects low frequency observations.
  • Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission: primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band.
  • Losses from the 3-bit samplers. The VLA's 3-bit samplers incur an additional 10–15% loss in sensitivity above that expected—i.e., the efficiency factor ηc = 0.78 to 0.83.

 

Table 3.2.1: SEFDs and D-Configuration Confusion Limits
Frequency   SEFD
(Jy)
 

RMS confusion level
in D config (µJy/beam)
Robust Weighting

0.39 GHz (P) 2790 5330
1.5 GHz (L) 420 74
3.0 GHz (S) 370 12
6.0 GHz (C) 310 2
10.0 GHz (X) 250 negligible
15 GHz (Ku) 320 negligible
20 GHz (K) 500 negligible
33 GHz (Ka) 600 negligible
45 GHz (Q) 1300 negligible

 

In general, the zenith atmospheric opacity to microwave radiation is very low: typically less than 0.01 at L, C, and X-bands; 0.05 to 0.2 at K-band; and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Observing conditions are best at night and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable.

Observers should remember that clouds, especially clouds with large water droplets (thunderstorms), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz.

Tipping scans—which are currently unavailable but will be implemented at some time in the future—can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the complex gain calibrator (amplitude and phase) and target source are observed.

When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following items:

  • Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target.

    Both CASA and AIPS allow the application of elevation-dependent gains and an estimated opacity generated from ground-based weather through the CASA tasks gencal and plotweather, and AIPS task INDXR.

  • Pointing. The SEFD quoted above assumes good pointing. Under calm, nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime can be much worse—occasionally exceeding 1 arcminte due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 11 mph (5 m/s). Wind speeds near the stow limit 45 mph (20 m/s) will have a similar negative effect at 8 and 15 GHz.

To achieve increased pointing accuracy, referenced pointing is recommended where a nearby calibrator is observed in interferometric pointing mode every hour or so. The local pointing corrections measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2–3 arcseconds (but more typically 5–7 arcseconds) if the reference source is within about 15 degrees in azimuth and elevation of the target source and the source elevation is less than 70 degrees. At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup.

Use of referenced pointing is highly recommended for all Ku, K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.3 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution.

To aid VLA proposers there is an online guide to the exposure calculator; the exposure calculator provides a graphical user interface to these equations.

Special caveats apply for P-band (230–470 MHz) observing. The SEFD's in Figure 3.2.2 or that listed in table 3.1.2 are from an observation taken far from the Galactic plane, where the sky brightness is about 30K. At P-band, Galactic synchrotron emission is very bright in directions near the Galactic plane. The system temperature increase due to Galactic emission will degrade sensitivity by factors of two to three for observations in the plane, and by a factor of five or more at or near the Galactic center. Additionally, the antenna efficiency (currently about 0.31 for 300 MHz) will decline with both increasing and decreasing frequencies from the center of P-band.

The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by:

where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼1.5λ/Bmax), and S in mJy per beam, the parameter F depends on the synthesized beam, therefore on the array configuration, and has the approximate value of F = 190, 18, 1.7, 0.16 for A, B, C, and D configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ΔIm, for S. Note that Equation 2 is a beam-averaged surface brightness; if a source size can be measured, then the source size and integrated flux density should be used in Equation 2 and the appropriate value of F calculated. In general, the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 6 of Reference 1, listed in Documentation.

For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791):

where D is the distance to the galaxy in Mpc, and SΔV is the HI line area in units of Jy km/s.

 

VLA Frequency Bands and Tunability

Bands

For observations taken with the 8-bit samplers, each receiver can tune to two different frequencies, each 1024 MHz wide, within the same frequency band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies, except for the low-band receiver (50–500 MHz), which provides linear polarization (X and Y). Each of these four data streams follows the VLA nomenclature and are known as IF (for Intermediate Frequency channel) A, B, C, and D. IFs A and B provide RCP (or Y when applicable), IFs C and D provide LCP (or X when applicable). IFs A and C are always at the same frequency, as are IFs B and D (but note that the A and C IFs frequency is usually different from the B and D frequency). We normally refer to these two independent data streams as IF pairs, i.e., the A/C pair and the B/D pair. In 8-bit mode, a maximum of 1024 MHz can be correlated for each IF pair (see the WIDAR Section), for a total maximum bandwidth of 2048 MHz. To distinguish this 8-bit system from the 3-bit system, these IF pairs are denoted A0/C0 and B0/D0.

More options are available with the 3-bit samplers. This system provides four (R,L) polarization pairs, each 2048 MHz wide. The A/C IF pair provides two sampled pairs, labelled A1/C1 and A2/C2, and the B/D IF pair provides two sampled pairs, labelled B1/D1 and B2/D2.

For more details on the 8-bit and 3-bit samplers see the VLA Samplers section.

The tuning ranges, along with default frequencies for continuum applications, are given in Table 3.3.1 below.

Table 3.3.1: Default frequencies for continuum applications
Band Range1 8-bit continuum applications (GHz) 3-bit continuum applications (GHz)
  (GHz) IF pair A0/C0 IF pair B0/D0 IF pair A1/C1 IF pair A2/C2 IF pair B1/D1 IF pair B2/D2
4 m (4) 0.054 – 0.0842 .070 – .082
90 cm (P) 0.20 – 0.503 0.224 – 0.4803
4 m (4)/90 cm (P) 0.054 – 0.0842/
0.20 – 0.503
0.224 – 0.4803 .070 – .082
20 cm (L) 1.0 – 2.04 1.0 – 1.54 1.5 – 2.04
13 cm (S) 2.0 – 4.0 2.0 – 3.0 3.0 – 4.0
6 cm (C) 4.0 – 8.0 4.5 – 5.5 5.5 – 6.5 4.0 – 6.0 6.0 – 8.0
3 cm (X) 8.0 – 12.0 8.0 – 9.0 9.0 – 10.0 8.0 – 10.0 10.0 – 12.0
2 cm (Ku) 12.0 – 18.0 13.0 – 14.0 14.0 – 15.0 12.0 – 14.0 14.0 – 16.0 16.0 – 18.0
1.3 cm (K) 18.0 – 26.5 20.2 – 21.2 21.2 – 22.2 22.0 – 24.0 24.0 – 26.0 18.0 – 20.0 20.0 – 22.0
1 cm (Ka) 26.5 – 40.0 32.0 – 33.0 31.0 – 32.0 33.0 – 35.0 35.0 – 37.0 29.0 – 31.0 31.0 – 33.0
0.7 cm (Q) 40.0 – 50.0 40.0 – 41.0 41.0 – 42.0 44.0 – 46.0 46.0 – 48.0 40.0 – 42.0 42.0 – 44.0

Notes:

1.  Listed here are the nominal band edges. For all bands, the receivers can be tuned to frequencies outside this range, but at the cost of diminished performance. Contact the NRAO Helpdesk for further information.
2. The 4-band system is currently under development. The default frequency range maximizes sensitivity of the system and provides a nominal bandwidth of 12 MHz and a channel resolution of 64 kHz.
3. The default setup for P-band will provide 16 subbands from the A0/C0 IF pair, each 16 MHz wide, to cover the frequency range 224–480 MHz. The channel resolution is 125 kHz. 
4. The default frequency setup for L-band comprises two 512 MHz IF pairs (each comprising 8 contiguous subbands of 64 MHz) to cover the entire 1–2 GHz of the L-band receiver.

 

Tuning Restrictions

In general, for all frequency bands except Ka, if the total span of the two independent IF pairs of the 8-bit system (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For K, Ka, and Q-bands—the only bands where a span greater than 8 GHz is possible—there are special rules:

  • At Ka-band, the low frequency edge of the A0/C0 IF pair must be greater than 32.0 GHz. There is no restriction on the B0/D0 frequency, unless the B0/D0 band overlaps the A0/C0 band when the latter is tuned at or near the 32.0 GHz limit. In this case, the Observation Preparation Tool (OPT) may not allow the requested frequency setups. Users wanting to use such a frequency setup are encouraged to contact the NRAO Helpdesk for possible tuning options.
  • At K and Q-bands, if the frequency span is greater than 8.0 GHz, the B0/D0 frequency must be lower than the A0/C0 frequency.

For the 3-bit system, the maximum frequency span permitted for the A1/C1 and A2/C2 IF pairs is about 5000 MHz. The same restriction applies to B1/D1 and B2/D2. The tuning restrictions given above for the separation and location of the 8-bit pairs A0/C0 and B0/D0 also apply to the 3-bit pairs, with A0/C0 replaced by A1/C1 and A2/C2, and B0/D0 replaced by B1/D1 and B2/D2.

 

VLA Samplers

The VLA is equipped with two different types of samplers, 8-bit with 1GHz bandwidth, and 3-bit with 2GHz bandwidth. The choice depends on your science goals and on technicalities described below.

The 8-bit Set consists of four 8-bit samplers running at 2.048 GSamp/sec. The four samplers are arranged in two pairs, each pair providing 1024 MHz bandwidth in both polarizations. The two pairs are denoted A0/C0 and B0/D0. Taken together, the four samplers offer a maximum of 2048 MHz coverage with full polarization. The frequency spans sampled by the two pairs need not be adjacent. Some restrictions apply, depending on band, as described in the section on Frequency Bands and Tunability.

The 3-bit Set consists of eight 3-bit samplers running at 4.096 GSamp/sec. The eight samplers are arranged as four pairs, each pair providing 2048 MHz bandwidth in both polarizations. Two of these pairs, denoted A1/C1 and A2/C2 cannot span more than 5000 MHz (lower edge of one to the higher edge of the other). The same limitation applies to the second pair, denoted B1/D1 and B2/D2. The tuning restrictions are described in the section on Frequency Bands and Tunability. Taken together, the eight 3-bit samplers offer a maximum of 8192 MHz coverage with full polarization.

 

Which set to use?

  • S, L, and 4/P-band observations, whether line or continuum, should use the 8-bit sampler set.
  • C and X-band continuum observations should use 3-bit samplers in order to exploit the full 4 GHz bandwidth: in spite of the 15% reduction in sensitivity that comes with 3-bit (at equal bandwidth to the 8-bit samplers—see below for details) and the reduced effective bandwidth after removing RFI, this still provides superior overall sensitivity. For more details we refer to EVLA memo 166.
      • Note: C-band is impacted by strong RFI caused by microwave links near 6 GHz in the A and B configurations. As a result, 3-bit data obtained with the standard setup are corrupted. We advise observers to use mixed 3-bit and 8-bit samplers. For more details, refer to the VLA Observing Guide.
  • Ku, K, Ka, and Q-band continuum observations should use the 3-bit samplers for maximum bandwidth.
  • Wide-band spectral line searches requiring more than 2 GHz span should use the 3-bit samplers.
  • Spectral-line observations which fit within two, possibly disjoint, 1 GHz bands should use the 8-bit set.
  • Simultaneous continuum and high resolution spectral line observation can use mixed 3-bit and 8-bit samplers. The 3-bit samplers in this case will be set up to deliver the continuum data, while the 8-bit samplers will be for the spectral line data. This mix mode can be used in C-band and higher.


Major Characteristics of each Set

The 8-bit samplers are warranted for observations at 4/P, L, and S-bands. The full analog bandwidth from the receivers fits within the 2048 MHz span covered by the samplers.

For the 3-bit samplers, users need to be aware of the following issues:

  • Sensitivity: compared to the 8-bit system, the sensitivity of the 3-bit samplers is worse by ~15% (at equal bandwidth). Alternatively, a given continuum noise level requiring on-source integration time T with the 8-bit (two bands of 1GHz), requires 0.33T with the 3-bit (4 bands of 2GHz, assuming the bandwidth is available from the front end).
  • Resonances: each of the eight 3-bit samplers on an antenna has a resonance about 3 MHz wide. Each resonance is independent of all others, so there is no correlated signal between antennas. The resonance degrades the spectrum in its narrow frequency range, but has little effect on continuum observing. Bandpass solutions will be affected, but can be interpolated over. Spectral-line calibration and images at the affected frequencies will show significant loss in sensitivity. The resonances are easily seen in autocorrelation spectra, and it is recommended that users, especially spectral-line users, utilize these to locate the compromised frequencies.
  • Amplitude Calibration: The traditional method for both 8- and 3-bit systems is to observe a flux-density calibrator, use self-cal to determine the antenna amplitude calibration factors (gains), and transfer the gains to the phase calibrator and target. For 3-bit samplers this procedure gives results good to 5% between elevations of 20–70 degrees. (Expect worse at the upper edge of Q-band and/or during bad weather.) The switched power data can be used to correct for system gain variations and works well for the 8-bit samplers. For 3-bit samplers, the Pdif depends on the Psum, i.e., Pdif is non-linear and its application will bias the resulting visibilities by 5–10%. The origin of this effect is understood, but we have not yet determined how best to compensate for it. Because of this, we do not recommend use of the Psum and Pdif data to calibrate visibilities from the 3-bit samplers. We do, however, recommend that the requantizer gains in the switched power data be applied to remove gain changes. For more information about the switched power, Psum, and Pdif, see EVLA memo 145.

 

Setting up the 8-bit or 3-bit Samplers

Either set requires an initial scan for each individual LO (frequency) tuning, during which power levels are optimized.

For the 8-bit system, a dummy scan of 1 minute duration is sufficient for each tuning. This  is usually done while the antennas are slewing at the start of an observing file, as the pointing direction of the antennas is not critical.

For the 3-bit system, the requirements are more demanding, see the section on 3-bit setup within the Guide to Observing with the VLA. The minimum setup time is 1 minute for each tuning to adjust the power levels and bandpass slopes across the 2GHz samplers. These values are retained and applied if the tuning is re-encountered in the same observation. Additionally, every time the LO setup is changed—whether or not it is new (e.g., changing from 8-bit X-band reference pointing back to target)—a scan of 30 seconds is needed to reset the subband gains (requantizers) in the correlator. For better amplitude calibration at high frequencies, the 3-bit initial setup should be near the elevation of the target, so do it after the first 8-bit setup described above. For 3-bit observing without 8-bit (e.g., C or X-band without reference pointing), the power variation with elevation is small, so the 3-bit setup can be done at any elevation.

For settings that use a mix of 3-bit and 8-bit samplers, the guidelines to set up the 3-bit samplers should be followed.


Other issues

The overhead for setup of 3-bit samplers can eat into observing time, especially for projects with many different LO settings, and/or sources all over the sky accompanied by band change, reference pointing, and requantizer reset for each direction. The impact is most severe for short scheduling blocks.

Polarization testing conducted so far indicates no degradation of performance by using the 3-bit samplers.

Field of View

Primary Beam

The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is θPB = 42/νGHz for frequencies between 1 and 50 GHz (L- through Q-band). At P-band the approximate value is θPB = 50/νGHz. New precise measurements of the primary beam shape have been reported in EVLA Memo 195; these allow for the correction of the primary beam attenuation in wide-field images. Both AIPS and CASA (5.0 and later versions) have these new parameters incorporated.

With the wide-bandwidths of the VLA it is necessary to account for the variation of the primary beam with frequency in order to achieve high-dynamic range images. For this and other imaging details we refer to the Limitations on Imaging Performance section of the OSS.

To achieve good sensitivity with a single-pointing observation, observers should take care to ensure that their targeted patch of sky fits within the primary beam (θPB) corresponding to the highest frequency of their observing band. If that is not possible, multiple overlapping pointings can be used to construct images of larger regions of sky through a technique known as mosaicking. Guidelines for mosaicking with the VLA are given in the Guide to Observing with the VLA.

Note: The Largest Angular Scale (LAS) that can be imaged by the array is independent of the Primary Beam's field of view or the use of mosaicking to increase the field of view. A table of the band- and configuration-dependent LAS is presented in the Resolution section of this document.

 

Chromatic Aberration (Bandwidth Smearing)

The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array.

These effects can be parameterized by the product of the fractional bandwidth (Δν/ν0) with the source offset in synthesized beamwidths (θ0HPBW). Table 3.5.1 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter and should be used to determine how much spectral averaging can be tolerated when imaging a particular field.

Table 3.5.1: Reduction in Peak Response Due to Bandwidth Smearing
(Δν/ν0)*(θ0HPBW)   Peak   Width
0.0 1.00 1.00
0.50 0.95 1.05
0.75 0.90 1.11
1.0 0.80 1.25
2.0 0.50 2.00

Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). Δν/ν0 is the fractional bandwidth; θ0HPBW is the source offset from the phase tracking center in units of the synthesized beam.

Note: The VLA correlator supports frequency averaging for single subarray and non-OTF observations. Currently this capability is limited to an averaging by a factor of 2 or 4  and only for wide-band continuum science projects (appropriate for C-band through Q-band observations). Observers interested in this capability should consult the EVLA memo 199 to assess the suitability of the frequency averaging in the correlator for their observations, because the extent of the bandwidth smearing is heavily dependent on the frequency averaging factor, the array configuration, and the observing frequency.

 

Time-Averaging Loss

The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parametrized, except for observations at δ = 90°. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with Δν/ν0 replaced by ωeΔtint, where ωe is the Earth's angular rotation rate, and Δtint is the averaging interval.

For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 in Documentation considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 3.5.2, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θΔtint)2, where θ is the distance from the phase center and Δtint is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see the section on Time Resolution and Data Rates) in the A and B configurations.

Table 3.5.2: Averaging Time for a Given Amplitude Loss
Amplitude loss
Configuration   1.0%   5.0%   10.0%
A 2.1 4.8 6.7
B 6.8 15.0 21.0
C 21.0 48.0 67.0
D 68.0 150.0 210.0

Note: The averaging time (in seconds) results in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines.

 

Note: For both the chromatic aberration and the time-averaging loss, the issue is not a simple reduction in amplitude for sources far from the phase center, but a convolution to the extent that a point source far from the phase center will become resolved due to bandwidth and/or time smearing. Furthermore, the description given above for the bandwidth smearing is based on the assumption that the radiation is monochromatic to parameterize the smearing, and does not take into account the consequences of having wide-bandwidths as is the case for the VLA. Therefore, while proposing and planning for VLA observations, and depending on the objectives of the science and the location of the sources of interest within the field, including confusing sources which may be far outside the science field, the above noted guidelines need to be used to conservatively estimate the proper channel width and the correlator integration time in order to minimize the effects of the bandwidth smearing and the time smearing, respectively.

 

Non-Coplanar Baselines

The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the VLA with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at 90 cm and 20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations.

Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS task IMAGR and CASA task clean. Refer to the package help files for these tasks, or consult with the NRAO Helpdesk for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the W-projection method available in CASA (see EVLA Memo 67 for more details).

Time Resolution and Data Rates

The default integration times for the various array configurations and frequency bands are as follows:

Table 3.6.1: Default Integration Times
Configurations Observing
Bands
Default
integration time
A, B, C, D 4 P 2 seconds
A L S C X Ku K Ka Q 2 seconds
B L S C X Ku K Ka Q 3 seconds
C, D X Ku K Ka Q 3 seconds
C, D L S C 5 seconds

Observations with the 3-bit (wideband) samplers, when applicable, should use these integration times. Observations with the 8-bit samplers may use shorter integration times, but these must be requested and justified explicitly in the proposal, and obey the following restrictions:

Table 3.6.2: Minimum integration times and maximum data rates
Proposal type

Minimum integration time

Maximum data rate
General Observing (GO) 50 msec up to 60 MB/s (216 GB/hr)
Shared Risk Observing (SRO) 50 msec > 60 MB/s (216 GB/hour) and up to 100 MB/s (360 GB/hour)
Resident Shared Risk Observing (RSRO) < 50 msec > 100 MB/s (360 GB/hr)

Note that integration times as short as 5 msec and data rates as high as 300 MB/s can be supported for some observing, though any such observing is considered Resident Shared Risk Observing. For these short integration times and high data rates there will be limits on bandwidth and/or number of antennas involved in the observation. Those desiring to utilize such short integration times and high data rates should consult with NRAO staff.

The maximum recommended integration time for any VLA observing is 10 seconds.

Observers should bear in mind the data rate of the VLA when planning their observations. For Nant antennas and integration time Δt, the data rate is:

Data rate ~ 45 MB/sec × (Nchpol/16384) × Nant × (Nant − 1)/(27×26) / (Δt/1 sec)
~ 160 GB/hr × (Nchpol/16384) x Nant × (Nant − 1)/(27×26) / (Δt/1 sec)
~ 3.7 TB/day × (Nchpol/16384) × Nant × (Nant − 1)/(27×26) / (Δt/1 sec)

Here Nchpol is the sum over all subbands of spectral channels times polarization products:

Nchpol = Σi Nchan,i × Npolprod,i

where Nchan,i is the number of spectral channels in subband i, and Npolprod,i is the number of polarization products for subband i (1 for single polarization [RR or LL], 2 for dual polarization [RR and LL], 4 for full polarization products [RR, RL, LR, LL]). This formula, combined with the maximum data rates given above, imply that observations using the maximum number of channels currently available (16384) will be limited to minimum integration times of ~2 seconds for standard observations, and 0.8 seconds for shared risk observations.

We note that frequency averaging in the correlator will reduce the total number of channels. Therefore, the data rate and the data volume will be reduced by the same channel averaging factor. See the Chromatic Aberration section for more details on the frequency averaging in the correlator and to assess its impact on your science.

These data rates are challenging for transfer and analysis. Data may either be downloaded via ftp over the Internet, or shipped on hard drives for large data sets or for those with slow Internet connections (please review the data shipping policy). For users whose science permits, the Archive Access Tool allows some level of frequency averaging in order to decrease data set sizes before ftp; note that the full spectral resolution will be retained in the NRAO archive for all observations.

Note: The data rate formula given above does not account for the auto-correlations delivered by WIDAR. Precise data rate values can be obtained through the use of the Resource Catalog Tool for proposing (RCT-proposing).

Radio Frequency Interference

The very wide bandwidths of the VLA mean that radio frequency interference (RFI) will be present in a far larger fraction of current observations than in observations made with the old systems. Considerable effort has gone into making the VLA's electronics as linear as possible, so that the effects of any RFI will remain limited to the actual frequencies at which the RFI exists. Non-linear effects, such as receiver saturation, should occur only for those very unlikely, and usually very brief, times when the emitter is within the antenna primary beam.

RFI is primarily a problem within the low frequency (C, S, L, and the low-band system) bands, and is most serious in the D configuration. With increasing frequency and increasing resolution comes an increasing fringe rate, which is often very effective in reducing interference to tolerable levels.

The bands within the tuning range of the VLA which are protected for radio astronomy are: 73.0-74.6 MHz, 322.0-328.6 MHz, 1400–1427 MHz, 1660.6–1670.0 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No significant external interference should occur within these bands.

VLA staff periodically observes the entire radio spectrum with the VLA, from 1.0 through 50.0 GHz with 125 kHz channel resolution, to monitor the ever-changing RFI spectrum. Users concerned about the precise frequencies of strong RFI, and the likelihood of being affected, are encouraged to peruse these plots. To access these plots, or for more information on RFI, including the impact of satellite transmissions, please see the RFI section in the Guide to Observing with the VLA.

Subarrays

The continuum subarray option offers two 1 GHz baseband pairs with the 8-bit samplers in up to 3 subarrays or four 2 GHz baseband pairs with the 3-bit samplers, with the same spectral channel and polarization product options as are available for wideband observing. The setup for each subarray is completely independent in terms of observing frequency, polarization products, and integration times.

When using three subarrays, there are some restrictions on the number of antennas in each subarray. The Baseline Board in the correlator treats each set of 4 antennas independently, using a separate column of correlator chips. With 8 such columns, the correlator can handle up to 8×4 = 32 antennas. The correlator configuration software requires that a given column not be split across subarrays. This does not matter when using only two subarrays, but forces some subtle restrictions when using three. For instance, one cannot observe with 9 antennas in each of 3 subarrays, because 9 antennas requires three columns (two with 4 antennas each, and one with 1 antenna); three subarrays of 9 antennas each would require 3×3 = 9 columns, one more than are actually available. Splitting the array into 10, 9, and 8 antennas is allowed, since the first two subarrays use 3 columns each, while the third uses only two.

Table 3.8.1 gives four examples of how correlator resources can be split into multiple subarrays. Antennas in each subarray are color-coded: red for subarray 1, green for subarray 2 (if present), and blue for subarray 3 (if present). The last column gives the number of antennas in each subarray (e.g., in the setup shown in the first row, subarray 1 has 10 antennas, subarray 2 has 9 antennas, and subarray 3 has 8 antennas). In all cases a total of 27 antennas are used. The columns are numbered in reverse order (C7 to C0) to match the numbering scheme used on the actual Baseline Boards.

Table 3.8.1: Some Possible Subarray Options
Number of antennas correlated using each Baseline Board column Number
of antennas
C7 C6 C5 C4 C3 C2 C1 C0
4 4 2 4 4 1 4 4 10 + 9 + 8
4 4 4 2 4 4 4 1 14 + 13
4 4 3 4 4 3 3 2 11 + 11 + 5
4 4 4 4 4 4 3 27

For more information on subarrays, please see the Subarray section in the Guide to Observing with the VLA.

Positional Accuracy & Astrometry

The position of a target can be determined to a small fraction of the synthesized beam, limited by atmospheric phase stability, the proximity of an astrometric calibrator, the calibrator-source cycle time, and the SNR on target.

In preparation for observing, the a priori position must be known to within the antenna primary beam, except perhaps for mosaicking observations. In the special case of using the phased VLA as a VLBI element, the a priori position must be accurate to within the synthesized beam of the array.

In post-processing, target positions are typically determined from an image made after phase calibration, i.e., correcting the antenna and atmospheric phases as determined on the reference source. The accuracy of the calibration determines the accuracy of the positions in the image. Note that phase self-calibration imposes the assumed position of the model, i.e., makes the position indeterminate. Therefore, an absolute position cannot be determined after self-calibration, but relative positions between features within a self-calibrated image are valid.

It may help to think of astrometry as two methods, narrow-field and wide-field.

Narrow-field astrometry

In narrow-field astrometry, the target is close to the phase tracking center and the antennas nod every few minutes between the target and a calibrator.  If no special calibration provisions are taken, under typical conditions, an astrometric accuracy of ~10% of the synthesized beam FWHM can often be obtained.  For example, an observation in Ka-band (~33 GHz) in A-configuration might reach an astrometric accuracy of ~10 milliarcseconds (mas).  When care is taken (special calibrations and ideal observing conditions), the accuracy can approach 1–2% of the synthesized beam, with a floor of ~2 mas.  If such accuracies are needed, we strongly recommend obtaining advice from VLA staff in setting up the observations.

Astrometric calibrators are marked J2000 A in the VLA calibrator list, and have a positional accuracy of ~2 mas. Other catalogs from the USNO and the VLBA are also useful, but offsets may exist between the VLA and VLBA centroids arising from extended structure in the particular source and the different resolutions of the arrays.

For studies of proper motion and parallax, the absolute accuracy of a calibrator may be less important than its stability over time. Close, or in-beam calibrators with poor a priori positions, can be used and tied to the ICRF reference frame in the same or separate observations.

Phase stability can be assessed in real time from the Atmospheric Phase Interferometer (API) at the VLA site, which uses observations of a geostationary satellite at ~12 GHz. Dynamic scheduling uses the API data to run a project under suitable conditions specified by the user. VLBI projects using the phased VLA will typically be fixed date and not dynamically scheduled.

Wide-field astrometry

Wide-field astrometry is used to determine the positions of targets within the primary beam, referenced to a calibrator within the beam or close by. In addition to the previous effects, there are distortions as a function of position in the field, from small errors in the Earth Orientation Parameters (EOP) used at correlation time, differential aberration, and phase gradients across the primary beam. With no special effort, the errors build up to roughly one synthesized beam at a separation of ~104 beams from the phase tracking center. Not all these errors are fully understood, and accurate recovery of positions over the full primary beam in the wide-band, wide-field case is a research area. These effects are handled somewhat differently in the post-processing packages. Check with VLA staff for more details via the NRAO Helpdesk.

Limitations on Imaging Performance

Imaging Fidelity

Image fidelity is a measure of the accuracy of the reconstructed sky brightness distribution. A related metric, dynamic range, is a measure of the degree to which imaging artifacts around strong sources are suppressed, which in turn implies a higher fidelity of the on-source reconstruction.

With conventional external calibration methods, even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is most often the effective phase stability of the telescope due to atmospheric/ionospheric fluctuations, although pointing errors and changes in atmospheric opacity may also be a limiting factor. If a good model of the sky brightness distribution exist (e.g., use of compact structures of sufficient strength—though a good model of resolved sources in the field of view may also be used), standard self-calibration can be counted on to improve the images. At low frequencies, where the dominant phase error is due to ionospheric plasma density fluctuations, more advanced techniques may be required to account for change of ionospheric phase across the field of view. Depending on the underlying nature of the errors, dynamic ranges in the thousands to hundreds of thousands can be achieved using these techniques. With the WIDAR correlator and wide-band receivers, i.e. large bandwidths resulting in high sensitivities, self-calibration methods can be extended to observations of sources with much lower flux densities than ever possible with the old VLA.

The choice of image reconstruction algorithm also affects the correctness of the on-source brightness distribution. The CLEAN algorithm is most appropriate for predominantly point-source dominated fields. Extended structure is better reconstructed with multi-resolution and multi-scale algorithms. For high dynamic ranges with wide bandwidths, algorithms that model the sky spectrum as well as the average intensity can yield more accurate reconstructions.

 

Invisible Structures at Large Angular Scales

An interferometric array acts as a spatial filter so that, for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Image reconstruction algorithms such as multi-resolution and multi-scale CLEAN may help reduce these negative bowls, but care must be taken in choosing appropriate scale sizes to work with.

Table 3.1.1 in the Resolution section gives the largest scale visible to each configuration/band combination.

 

Poorly Sampled Fourier Plane

Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in Documentation.

 

Sidelobes from non-Deconvolved Sources

At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (tclean) are well suited to this task. Sidelobe confusion is a strong function of observing band—affecting most strongly L and P-band observations—and is rarely a significant problem for observations at frequencies above 4 GHz.

 

Sidelobes from Strong Sources

An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral-line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere.

 

Wide-band Imaging

The very wide bandpasses provided by the VLA enable imaging over 2:1 bandwidth ratios: at L, S, and C-bands, the upper frequency is twice that of the lower frequency. It is this wide bandwidth which enables sub-microJy sensitivity.

In many cases, where the observation goal is a simple detection and there are no strong sources near to the region of interest, standard imaging methods that combine the data from all frequencies into one single image (multi-frequency-synthesis) may suffice. A rough rule of thumb is that—provided a strong source is not adjacent to the target zone—if the necessary dynamic range in the image is less than 1000:1 (i.e., the strongest source in the beam is less than 1000 times higher than the noise), a simple wide-band map may suffice.  

For higher dynamic ranges, complications arise from the fact that the brightness in the field of view dramatically changes as a function of frequency, both due to differing structures in the actual sources in the field of view and due to the attenuation of the sources by the primary beam. One symptom of such problems is the appearance of radial spokes around bright sources, visible above the noise floor, when imaged as described above.  

The simplest solution is to simply make a number of maps (e.g., one for each subband) which can then be suitably combined after correction for the primary beam shape. But with up to 64 subbands available with the VLA's new correlator, this is not always the optimal approach, both from a practical standpoint and frequency-dependent image reconstruction uncertainties. Further, images at all bands must be smoothed to the angular resolution at the lowest frequency before any spectral information can be extracted. With a 2:1 bandwidth, the difference in angular resolution across the band will be significant.

A better approach is to process all subbands simultaneously, utilizing software which takes into account the possibility of spatially variant spectral index and curvature and knows the instrumentally-imposed attenuation due to the primary beam. Such wideband (plus wide-field) imaging algorithms are now available within CASA as part of the tclean task.

 

Wide-field Imaging

Wide-field observing refers to both the non-coplanar nature of the VLA when observing in non-snapshot mode and the effect of the primary beam patterns.

At high angular resolutions and low frequencies, standard imaging methods will produce artifacts around sources away from the phase center due to the non-coplanar nature of the VLA. Faceted imaging (AIPS, CASA) and W-projection (CASA) techniques can be used to solve this problem. Community image packages also employ similar techniques, usually geared toward addressing low frequency and wide-field imaging problems and might work with VLA data as well. 

Wide field imaging also requires the accurate representation of primary beam patterns and their use during imaging. This is relevant only for very high dynamic ranges ( > 10000:1 ) or when there are very strong, confusing sources at and beyond the half-power point of the primary beam. This problem is worse with a wide-band instrument simply because the size of the primary beam—and the radius at which the half-power point occurs—varies with frequency while there is increased sensitivity out to a wider field of view. In CASA, the tclean task contains algorithms to deal with these effects by modeling and correcting for polarization squint, frequency-dependent and rotating primary beams, per antenna, during imaging. Work is under way to develop full-polarization primary beam corrections.  Please note, however, that most advanced methods are computationally very demanding, will lead to a significant increase in processing time, and may not always be required. Therefore, in the interest of practicality, they should be used only if there is evidence of artifacts without these methods. At low frequencies, wide-field imaging typically includes both non-coplanar effects plus primary beams whereas at high frequencies it is primarily primary beam effects.

Mosaicking is another form of wide-field imaging in which data from multiple pointings are combined either during imaging as a joint mosaic (CASA) or after imaging by stitching several images together (CASA, AIPS). It is also possible to apply a- and w-term corrections for joint mosaicking in CASA, however in most situations this requires significant computational resources and cannot be provided through standard NRAO post-processing computing. Furthermore, the current implementation in tclean requires for this case particular attention to run efficiently in a cluster environment. For more and up-to-date information on this topic, please refer to the CASA documentation on synthesis imaging.

 

Calibrating the Flux Density Scale

Normal calibration of the flux density scale for VLA observations is effected by including a scan on a source of presumed known flux density in each Scheduling Block (SB). Using that known flux density source, the flux density of the complex gain calibrator(s) can be determined and then transferred to your target source(s). Historically, 3C48 and 3C286 have been the standard sources for which NRAO has assumed flux densities are known as a function of frequency, and which have been recommended as flux density scale calibrator sources for the VLA. Restrictions on baseline length as a function of VLA configuration and observing band were supplied which, if followed, allowed relatively accurate flux density scale calibration. We have recently improved the ability to calibrate the flux density scale by providing sky brightness models for these sources in CASA and AIPS, which loosens the restrictions on configurations and bands. We have also added the sources 3C138** and 3C147 to the list of calibrators that have models. However, 3C48*, 3C138**, and 3C147 have spectral flux densities that vary with time (3C286, along with 3C295 and 3C196, are constant), so some care should be taken if the most accurate flux density scale calibration is desired.

Note: While accurate models are available in both AIPS and CASA for various frequency bands for the calibrators 3C286, 3C48*, 3C147, and 3C138**, neither 3C295 nor 3C196 has such models in CASA. Therefore, the VLA CASA calibration pipeline will fail if these two calibrators are used. Furthermore, 3C295 and 3C196 may not be suitable for all VLA configurations and frequencies even if one chooses to not use the pipeline.

A single observation of a few minutes of one of the above-mentioned flux density scale calibrators will suffice for most observers. If possible, the flux density scale calibrator should be observed at a time when it is nearly at the same elevation as the complex gain calibrator, especially for the highest four bands (Ku-, K-, Ka-, and Q-band). This is not always possible because of timing and geometry of sources, and that it is not typically known when an SB will be executed (so elevations versus time are uncertain). Flux density scale calibration accuracy in this case should be of order 10% at 4- and P-bands, 5% at L- through Ku-bands, and 10-15% for the three higher bands. If more accuracy is needed, a more careful strategy should be adopted, potentially using multiple flux density scale calibrators. The fundamental accuracy of the scale is ~5% at 4- and P-bands, 3% at L- through Ku-bands, increasing to 5% at Q-band. See Perley and Butler (2017) for more details on how the spectral flux densities of 3C48*, 3C138**, 3C147, and 3C286 (and many other sources) have been determined across the frequency range from 50 MHz to 50 GHz and how they vary versus time, along with information on the fundamental accuracy of the flux density scale when using these sources.

If less accuracy is needed in the flux density scale calibration, an observation of one of these standard sources need not necessarily be included in an SB. As an example, for a short triggered observation where a simple detection is desired, the time spent slewing back and forth to the flux density scale calibrator can make the SB significantly longer than it could otherwise be. In this scenario, the switched power measurement can be used to calibrate the flux density scale; see EVLA Memo 120 for some background. This technique, which should only be used with the 8-bit samplers, is not a standard path of calibration, but it is possible. The flux density scale accuracy in this case is ~10% for L- through Ku-bands, increasing to ~20% at Q-band; not nearly as good as using the "standard" method of flux density scale calibration, but it may be sufficient for some observers.

For reference, the polynomial expression for the spectral flux density for 3C286 determined in Perley and Butler (2017) is: \[\log(S) = 1.2481 - 0.4507 \log(f) - 0.1798 \log^2(f) + 0.0357 \log^3(f)\] where S is the flux density in Jy, and f is the frequency in GHz.

The tables below show flux densities determined using the polynomial coefficients for a few sources at a single frequency within each of the VLA bands.

Flux densities (Jy) of Standard Calibrators for January 2016
Source 75 MHz 350 MHz 1500 MHz 3000 MHz 6000 MHz 10000 MHz 15000 MHz 22000 MHz 33000 MHz 45000 MHz
3C48* = J0137+3309 72.8 42.2 15.4 8.44 4.42 2.68 1.79 1.22 0.815 0.601
3C138** = J0521+1638 26.5 16.1 8.25 5.44 3.39 2.33 1.72 1.28 0.949 0.761
3C147 = J0542+4951 58.0 52.3 21.0 12.0 6.45 3.99 2.73 1.93 1.39 1.13
3C196 = J0813+4813 129 44.4 13.6 6.98 3.38 1.91 1.20 0.763 0.473 0.329
3C286 = J1331+3030 30.0 25.9 14.6 9.91 6.39 4.50 3.37 2.54 1.88 1.49
3C295 = J1411+5212 124 58.4 21.2 11.0 5.06 2.70 1.60 0.970 0.571 0.385
Flux densities (Jy) of Standard Calibrators for January 2019
Source 328 MHz 1465 MHz 2565 MHz 4885 MHz 6680 MHz 11320 MHz 16564 MHz 25564 MHz 32064 MHz 48064 MHz
3C48* = J0137+3309 43.9 15.6 9.82 5.48 4.12 2.56 1.86 1.33 1.11 0.816
3C138** = J0521+1638 15.9 8.26 6.00 4.00 3.23 2.24 1.69 1.25 1.06 0.821
3C147 = J0542+4951 53.9 21.4 13.8 7.88 5.91 3.67 2.61 1.82 1.53 1.14
3C196 = J0813+4813 46.5 13.8 8.14 4.22 3.00 1.67 1.08 0.656 0.508 0.313
3C286 = J1331+3030 25.8 14.6 10.9 7.33 5.97 4.12 3.15 2.30 1.92 1.44
3C295 = J1411+5212 61.1 21.6 12.8 6.42 4.45 2.33 1.43 0.819 0.596 0.405

We refer the reader to the VLA Observing Guide for the practical considerations (e.g., observing frequency and array configuration, as well as post processing) regarding the choice of the flux density scale calibrator in their scheduling blocks.

* The flux density scale calibrator 3C48 has been undergoing a flare since January 2018 or so.  While we have not fully characterized this with the VLA, other instruments have measured it at some frequencies. At Ku-band the magnitude of the flare is of order 10%.  The effect will be smaller at lower frequencies (of order 5% at L-band), and might be larger at higher frequencies (of order 20% at Q-band).  If you care about the flux density scale of your observations at that level, you may want to re-calibrate your data once new time-variable values have been put into CASA and AIPS.

** The flux density scale calibrator 3C138 is currently undergoing a flare. From VLA calibration pipeline results, we have noticed that 3C138 is deviating from the model. The amount of this deviation is still being investigated by NRAO staff, but does seem to effect frequencies of 10 GHz and higher. At K and Ka-bands the magnitude of the flare is currently of order 40-50% compared to Perley-Butler 2017 flux scale. If you care about the flux density scale of your observations above 10 GHz, monitoring datasets are publicly available in the archive under project code TCAL0009, from which you may find an updated flux density ratio to use for your data.

Complex Gain Calibration

General Guidelines for Complex Gain Calibration

Adequate complex gain calibration (tracking amplitude and phase fluctuations as a function of time) is a complicated function of source-calibrator separation, frequency, array scale (configuration), and weather. Since what defines adequate for some experiments is completely inadequate for others, it is difficult to define simple guidelines to ensure adequate phase calibration. However, some general statements remain valid most of the time. These are given below.

  • Under decent conditions with no thunderstorms or ionospheric storms, tropospheric effects dominate at frequencies higher than about 4 GHz; ionospheric effects dominate at frequencies lower than about 4 GHz.
  • Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the C and D configurations at L and S-bands, and in the D configuration at X and C-bands. For these cases, calibration need only be done to track instrumental changes—three or four times per hour is usually sufficient for tracking the system gains.
  • If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than a couple of times per hour at low frequencies or 15 minutes at high frequencies in order to track pointing or other effects that might influence the amplitude scale. The enhanced sensitivity of the VLA guarantees, for full-band continuum observations, that every field will have enough background sources to enable phase self-calibration at L and S-bands. At higher frequencies, the background sky is not sufficient, and only the flux of the target source itself will be available.
  • In principle, the smaller the source-calibrator angular separation, the better (even if the closer calibrator is weaker). However, the final choice will depend on the observing frequency. If deciding between a nearby calibrator with an S code in the calibrator database, and a more distant calibrator with a P code, for low frequencies (L-band and below) the nearby calibrator is usually the better choice (low frequency strategy). For higher frequencies it is advisable to use the further away but higher quality P-code calibrator (high frequency strategy). A detailed description of calibrator codes is available in the calibrator list.
  • Phase stability often deteriorates dramatically after about 10AM due to small-scale convective cells set up by solar heating, even in clear and calm conditions, especially in the summer. Observers should consider a more rapid calibration cycle for observations between this time and a couple of hours after sundown.
  • At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is needed to track tropospheric phase fluctuations if the target cannot be self-calibrated. See Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) (below).

 

Rapid Phase Calibration and the Atmospheric Phase Interferometer

For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-calibrator observing cycles as short as 40 seconds can be used for very small source-calibrator separations. Observing efficiency declines, however, for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. This is represented in the Observation Preparation Tool as a loop of source-calibrator scans with short scan length. This technique stops tropospheric phase variations at an effective baseline length of ∼vat/2, where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec) and t is the total switching time. Short source-calibration scans have been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Note that at a typical wind velocity in the compact D-configuration, this effective baseline length is the same as—or larger than—the longest baseline in the array and it is not worth the increased overhead of short cycle times. Under these circumstances, it is sufficient to calibrate every 5-10 minutes to track the instrumental changes. The fast switching technique will not work in bad weather (such as rain showers or when there are well-developed convection cells (thunderstorms)). It is important to correctly specify the required tropospheric phase stability as measured by the Atmospheric Phase Interferometer at observe time (see below).

Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from References 9 and 10 in Documentation.  Also see the High Frequency Strategy guide for additional recommendations on observing at high frequencies.

An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase. The API uses an interferometer of two, 1.5 meter antennas, separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data are heavily used for the dynamic scheduling of the VLA.

Characteristic seasonal averages are shown in Table 3.12.1 below:

Table 3.12.1: Seasonal API/wind values at the VLA
Month

API (night)
[deg]

API (median)
[deg]

API (day)
[deg]

Wind (night)
[m/s]

Wind (median)
[m/s]

Wind (day)
[m/s]

January 2.3 2.8 3.6 1.6 1.9 2.3
February 2.9 3.4 4.5 4.0 4.3 4.5
March 2.8 3.7 5.5 3.4 3.9 4.7
April 3.3 4.5 6.2 5.3 5.5 5.8
May 2.9 4.6 6.7 2.6 3.2 3.7
June 3.8 5.5 7.4 2.5 3.9 6.3
July 6.2 8.3 10.5 2.9 2.9 3.0
August 5.4 7.1 11.3 1.7 2.3 3.0
September 5.2 6.6 8.8 2.3 3.0 3.6
October 4.2 5.3 7.4 2.3 2.9 3.7
November 2.6 3.0 4.0 1.2 2.5 1.6
December 2.8 3.2 4.1 1.2 1.6 2.7

Day indicates sunrise to sunset values; night indicates sunset to sunrise values.

 

Other Issues that Affect Complex Gain (Amplitude)

There are other instrumental effects that cause fluctuations in gain over time for VLA observations - we describe three of them here.  Note that all of these effects are to the gain amplitude, not phase.

Gain Curves

The VLA antennas have elevation-dependent gain variations which are important to account for at the four highest frequency bands. Gain curves are determined by VLA staff per antenna and per band, and the necessary corrections can be applied to the visibility data in both AIPS and CASA. Additionally, atmospheric opacity will also cause an elevation-dependent gain which is also particularly notable at these four highest frequency bands. We currently do not have an atmospheric opacity monitoring procedure; users should utilize the appropriate tasks available in both AIPS and CASA to estimate and correct for the opacity using ground-based weather data. If your complex gain calibrator is near your target source, and the flux density scale calibrator is also observed at a similar elevation, then most of this elevation-based gain will be calibrated correctly during normal calibration. Note also that a good procedure for removing elevation-based gain dependencies uses the AIPS task ELINT. This task will generate a 2nd order polynomial gain correction utilizing your own calibrator observations. This will remove both the antenna and opacity gain variations, and has the decided advantage of not utilizing opacity models or possibly incorrect antenna gain curves. Use of this procedure is only practical if your observations span a wide range in elevation.

Antenna Pointing Offsets

Another important gain variation effect at the four highest frequency bands is that due to antenna pointing offsets. Daytime observations on sunny days can suffer pointing errors, primarily in elevation, of up to one arcminute. This effect can be largely removed by utilizing the referenced pointing procedure which determines the pointing offset of a nearby calibrator. This offset is then applied to subsequent target source observations. It is recommended that this local offset be determined at least hourly, more often during sunrise or sunset, utilizing an object within 15 degrees of the target source—preferentially at an earlier right ascension. Studies show that the maximum pointing error will be reduced to about 7 arcseconds or better if proper referenced pointing is utilized.

Electronic Gain Variations

The VLA's post-amplifiers are not temperature stabilized and exhibit significant gain (amplitude) changes between night and day, particularly at the four highest frequency bands. Changes as large as 30% have been seen between night and day in calm, clear conditions. These gain changes, and others caused by possible changes in attenuator settings, are monitored and can be removed by application of the internal calibration signal, whose results are recorded in the switched power table in both AIPS and CASA. These corrections are not applied by the calibration pipeline—users who wish to correct for these gain changes must utilize the appropriate tasks in AIPS or CASA. For the most accurate flux density bootstrapping, this table must be applied to the visibility data before calibration.

    Polarization

    For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle or a single observation of an unpolarized source. The complex gain calibrator chosen for the observations can also double as a polarization calibrator, provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15–50 degrees during a several hour track).

    The minimum condition that will enable accurate polarization calibration from a polarized source, in particular with unknown polarization, is three observations of a bright source spanning at least 60 degrees in parallactic angle (schedule four scans, if possible, in case one is lost); if at all possible, it is strongly recommended that five or more observations covering 100 degrees (or more) of parallactic angle in roughly uniform steps be run. If a bright, unpolarized, unresolved source is available, and known to have very low polarization, then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for small objects as compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission; 3C48 can be used at frequencies of ~3 GHz and higher, or 3C147 at frequencies over ~10 GHz. Note that 3C48 and 3C138 are variable—the polarization properties are known to be changing significantly over time, most notably at the higher frequencies (for details see Perley and Butler (2013b)).

    More information on polarization calibration strategy can be found in the Polarimetry section in the Guide to Observing with the VLA.

     

    VLBI Observations

    The VLA can participate in VLBI observations with the VLBA. This is possible by either using the VLA as a phased array (Y27) or using one of its dishes (single dish: Y1). We note that currently P-band cannot be phased. For more details see the VLBI at the VLA documentation. In phased array mode, the program TelCal derives the antenna-based delay and phase corrections needed for antenna phasing in real time. This correction is applied to the antenna signals before they are summed, requantized to 2-bits, and recorded in VDIF format on the Mark5C disk at the VLA site. The disk(s) are then transported to Socorro, NM and correlated on the DiFX correlator with other VLBI stations which participated in the observation.

    Standard VLA data, i.e., correlations between VLA antennas, are also archived in the NRAO science data archive. By default, a maximum of 512 MHz dual polarization is phased/recorded and sent to the archive. This leaves a large portion of the possible VLA bandwidth unused (up to 2 or 8 GHz total depending on sampler choice). Increasing the bandwidth to the maximum that the VLA can provide has obvious benefits for most continuum observations. Pulsar gating can be performed commensally with the phased VLA, avoiding the need for extra single dish observations. Line VLBI observations that run through the WIDAR correlator currently produce each requested VLBI subband bandwidth divided into full polarization products with 64 frequency channels, regardless of the requested spectral resolution obtained in the VLBI correlation. By adding unused baseline boards to the VLBI-specified line subbands, a closer match in spectral resolution can be offered for the standalone VLA data. If you think your science could benefit from these capabilities, please review our documentation on VLBI at the VLA or contact the helpdesk for further guidance.

    Snapshots

    The two-dimensional geometry of the VLA allows a snapshot mode whereby short observations can be used to image relatively bright, unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest.

    Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm and longer wavelengths in the D configuration. For instance, at 20 cm, a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. The deconvolution of the data is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing Faint Images of the Radio Sky at Twenty-centimeters survey (FIRST, http://sundog.stsci.edu/top.html) (B configuration) or the NRAO VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) (D configuration, all-sky).

    Shadowing and Cross-Talk

    Observations at low elevation in the C and D configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling when using the default inputs with CASA tasks importasdm and importevla. AIPS task UVFLG can be used to flag VLA data based on shadowing, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other subarrays. The CASA task flagdata can also be used to flag data based on shadowing. For more information on shadowing, please see the Antenna Shadowing section in the Guide to Observing with the VLA.

    Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. This effect is important at low frequencies in compact configurations. Careful examination of the visibilities is necessary to identify and remove this form of interference. The affected data would show time-variable high-amplitude points.

    Combining Configurations and Mosaicking

    Any single VLA configuration will allow accurate imaging of a range of spatial scales determined by the shortest and longest baselines. For extended and structured objects, it may be required to obtain observations in multiple array configurations. It is advisable that the frequencies used be the same for all configurations to be combined. The ideal combination of arrays results in a uv-plane with all cells equally filled by uv-points. To first order, this can be achieved by using the beam sizes of the individual arrays to inversely scale the on-source integration time. This approach is equivalent to achieving the same surface brightness sensitivity for all arrays on all scales. For the VLA, observations in the different configurations generate beam sizes that decrease by factors of three, i.e., C configuration generates a three times smaller beam than D configuration, B three times smaller than C, and A three times smaller than B. Thus, on-source integrations would increase by about an order of magnitude between each array. Such a drastic increase is very expensive and, in fact, not necessary since some spatial scales are common to more than a single array, which is equivalent to some uv-cells being filled more than others. The best way to fill the uv-plane depends on many factors such as declination of the source, LST time of the observation, and bandwidth.

    Experience shows for the VLA that a factor of about three in on-source integration time for the different array configurations works well for most experiments. For example, a 20min on-source time in D, 1hr in C, 3hrs in B, and 9hrs in A should produce a decent map. Using large bandwidths and multi-frequency synthesis will broaden all uv tracks radially and one may need even fewer array configurations or shorter integration times between the different arrays.

    Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicking. The VLA has no limit on the number of pointings for each mosaic. Typically hexagonal, rectangular, or individual pointing patterns are used and the overlap regions will result in an improved rms over each individual pointing. Given the many, potentially short observations, it is important to obey the data rate limits outlined in the Time Resolution and Data Rates Section. In addition to discrete or pointed mosaics, on-the-fly (OTF) mosaics (i.e. dumping the data while moving the telescopes across the source) are also available.

    Time-variable structures, such as the nuclei of radio galaxies and quasars, cause special, but manageable, problems. See the article by Mark Holdaway in Reference 2 of the Documentation for more information.

    Guidelines for mosaicking with the VLA are given in the Guide to Observing with the VLA.

    Pulsar Observing

    The VLA can be used for several kinds of pulsar observing: phase-binning using the WIDAR correlator, using the phased-array for single-beam pulsar processing in either search or fold modes, or simply standard imaging mode with fast integrations. Both phase-binning and phased-array (YUPPI) modes are available under General Observing (GO). The only exception is the 4-band YUPPI which is a Resident Shared Risk Observing (RSRO) capability. For any questions not addressed here regarding the capabilities of these observing modes, please contact the NRAO Helpdesk.

    Phased-array pulsar processing

    The "Y" Ultimate Pulsar Processing Instrument (YUPPI) is a software suite that runs in the correlator backend (CBE) computer cluster and can process a single-beam phased/summed-array data stream for pulsar observations in real time, into either folded profiles or search mode (filterbank) output. Coherent dedispersion can be optionally applied in either mode.

    In the phased-array pulsar processing mode, the voltage data streams from each antenna are divided into a number of frequency subbands within the correlator, then summed and requantized before being output to the cluster for pulsar processing. The limitation on bandwidth comes primarily from the available network connections between the correlator and cluster. In all cases, a maximum of 64 subbands total can be processed. Depending on the number of bits chosen, this results in the following total bandwidth constraints:

    Table 3.18.1: Pulsar Observing Bandwidth Constraints

    Subband bandwidth

    Subband quantization

    Max total bandwidth

    Samplers

    32 MHz 8 bits 2048 MHz 8-bit
    64 MHz 4 bits 4096 MHz 3-bit
    128 MHz 2 bits 8192 MHz 3-bit
    <32 MHz 8 bits 64*BWsub 8-bit

    As described in the VLA Frequency Bands and Tunability section, the 8-bit samplers provide two independently tunable 1 GHz IFs, while the 3-bit samplers provide four tunable 2 GHz IFs. 

    The pulsar-specific processing is done in real time using the DSPSR software package and, in principle, any processing option supported by DSPSR can be used; this will be constrained by the real-time computing power available in the cluster. In general, each subband can be divided into an arbitrary (2n) number of channels; 1 (summed), 2 or 4 detected polarization products can be output; and coherent dedispersion can be enabled or not.

    Fold mode

    In fold mode, the data are averaged modulo a known pulsar ephemeris (provided via a standard TEMPO/TEMPO2 "par file") into pulse profiles. The data can also be folded at a constant topocentric period, for example at 10 Hz to detect the injected noise cal signal. Fold integration times as short as 1 second have been tested. Up to 16384 profile bins can be used. The data are recorded in PSRFITS format using the standard 16-bit data encoding. This means the final output data rate is given by:

    Data rate = 2 bytes × Nsubband × Nchannel × Nbin × Npoln / Tint

    If the desired data rate exceeds ~25 MB/s, additional testing ahead of time may be required.

    Search mode

    In search mode, the data are simply detected and averaged over a specified amount of time before being output to disk, resulting in a filterbank data array (power vs time and frequency). Coherent dedispersion at a known DM can optionally be enabled for this.  Data can be recorded using 2, 4, 8, 16 or 32 bits, resulting in a final data rate of:

    Data rate = (Nbit/8) bytes × Nsubband × Nchannel × Npoln / Tint

    The maximum sustained output rate in this mode should be kept less than ~400 MB/s.

    Subarrays

    It is possible to use any of the phased-array pulsar modes listed here in a subarray observation, following the guidelines described in the Subarrays section. In addition, the above constraints on the pulsar processing apply to the total of all simultaneously-used subarrays, rather than each subarray separately. For example, the total number of subbands in use across all subarrays must not be greater than 64; the total (not per-subarray) data rate must meet the above constraints, etc. It is possible to use different parameters such as subband bandwidth, number of bits, or processing mode (fold versus seach) in the different subarrays.

    VLBI

    It is possible to use phased-array pulsar processing as part of a VLBI experiment; see the VLBI Observations section, and links therein, for additional information about VLBI at the VLA. The main constraint on this type of observation is that a subband that is being recorded for VLBI can not be sent to the pulsar processing system. However, since VLBI recording typically only requires a small number of subbands (2 through 8), any additional subbands produced by WIDAR can be sent to the pulsar system, following the constraints above. This provides a high time resolution data stream covering wider bandwidth than the VLBI data. One typical use case is to detect a pulsar using the VLA data stream and determine a short-term timing ephemeris covering the observation. This can then be used to gate the VLBI correlation, reducing uncertainties associated with extrapolating existing timing solutions or obtaining time on other telescopes for these purposes.

    Gated or binned visibilities

    The WIDAR correlator has the capability to internally integrate (fold) visibilities into 1 or more pulse phase bins, following a standard TEMPO-compatible pulsar ephemeris. This mode can be used to image the emission from a pulsar of known period anywhere in the telescope field of view. This provides both higher signal-to-noise ratio on the pulsar than a standard image, and allows the pulsed emission to be separated from continuous emission from other sources in the field.

    The following constraints apply to binning-mode observations:

    • Binning is limited to the case where the pulse period is divided evenly into bins covering the full pulse period.  "Gating" style observations (common in VLBI) where a single on-pulse bin is used are not supported.
    • The maximum number of pulse phase bins is 1000.
    • There is a tradeoff between the total bandwidth and the minimum bin width (pulse period divided by number of bins):
      • With 4 subbands (up to 512 MHz total), the minimum allowed bin width is 12.5 μs.
      • With 16 subbands (up to 2048 MHz total), the minimum allowed bin width is 50 μs.
      • With 64 subbands (up to 8192 MHz total), the minimum allowed bin width is 200 μs.
    • The number of channels per subband is currently limited to 128 maximum.  Combining recirculation and binning is not allowed.
    • Integration (dump) time must be an integer number of pulse periods.
    • The data rate produced in this mode is the standard VLA data rate (see the Data Rate section) multiplied by the number of bins.  The data rate must be kept less than 60 MB/s.

    It should also be noted that there is currently very limited support for binned observations in standard data processing software (e.g., CASA). Development of data analysis procedures is ongoing and users of this mode should be aware that this will likely involve some advanced/low-level manipulation of raw VLA data sets.

    Fast-dump visibilities

    While not specifically a pulsar mode, standard visibility data can be dumped as fast as 5 ms, which may be sufficient for imaging of slow pulsars. See the Time Resolution and Data Rates section for more details.

    Solar Observing

    Observations of the Sun can currently be performed in five frequency bands: 1-2 GHz (L band), 2-4 GHz (S band), 4-8 GHz (C band), 8-12 GHz (X band), and 12-18 GHz (Ku band). To observe the Sun, any 8-bit resource using these bands can be used in Solar mode (see below). As the Sun is moving across the sky, the source position should be defined in the source catalog.

    If the center of the solar disk needs to be tracked, select the Sun using the drop-down for "Solar System Body with Internal Ephemeris". Otherwise, to allow for specifying the location of interest on or near the solar disk and the differential rotation model to be used for tracking the source, solar observations require selection of a source position type of "Solar System Body with Uploaded Ephemeris". An ephemeris file can be generated using the guidelines given in the Moving Object section or, for example, by using the ALMA Solar Ephemeris Generator

    The user must select the “Solar” scan mode when specifying scan details under Observation Preparation in the OPT. This scan mode ensures that the necessary attenuators are switched into the signal path when observing the Sun and that special noise cal sources are employed in the switched power system; this allows users to flux calibrate their data. Two observing modes are currently supported:

    • For quiet Sun observations, when no strong active regions or flares expected, users should select a scan mode of “Solar Attenuators with Low Noise Internal Cal”.
    • If the science objective involves observing a strong active region and/or flare activity, users should select a scan mode of “Solar Attenuators with High Noise Internal Cal”.

    These selections ensure that cal noise levels of these scans are of order a few percent of the anticipated system temperature. (A third scan mode, “Noise Reverse Coupler Setup (L Band only)”, is not yet supported.)

    Solar observations rely on the switched power system to flux calibrate the data. This, in turn, requires the use of 8-bit sampling modes in the frequency bands that support solar observing. Solar observing is such that input power levels may change from scan to scan, particularly for active Sun targets. This being the case each source scan requires a setup scan to trigger the system to reset the requantizer gain.

    Considerations for the calibration of solar observations are otherwise quite similar to standard gain calibration. Phase calibration proceeds in the same manner for solar observations as it does for general observing. However, it is important to be aware of two factors: First, when observing a calibrator source, the attenuation used when observing the Sun is removed from the signal path; the RFI that is suppressed, to a large degree, when observing the Sun is fully present in observations of calibrators. Second, when the attenuators are removed, solar radiation may be present in the sidelobes of the primary beam if the phase calibrator is too close to the Sun, particularly if the Sun is in an active phase. Therefore, it is advisable to observe phase calibrators that are further in angular distance from the source than would normally be used.

    VLA+LWA (eLWA) Observing

    The VLA can be used to record individual antenna signals (voltages) to VLBI Data Interchange Format (VDIF) files. Similar to Very Long Baseline Interformetry, using the phased-VLA or individual antennas, this specifically allows for joint correlation of VLA antennas with the Long Wavelength Array (LWA) stations in New Mexico in the overlapping 4m-band range. Joint correlation is performed offline using the LWA software correlator, which is located inside the VLA control building and produces FITS-IDI compatible data output.

    For shared risk observing, this mode is available for a single subband with a center frequency of 76 MHz, a bandwidth of 8 MHz, and 4-bit VDIF output. Other possible modes or center frequencies (limited by the VLA MJP-dipole response) are available through resident-shared risk observing. Use of any RSRO modes should also be brought to the attention of the LWA Director in order to verify that correlation is possible.

    eLWA proposals that are granted VLA time are automatically granted time for the LWA stations, for more information on how to select the eLWA resource for your proposal, refer to the PST manual. Observations are prepared and scheduled like any other VLA observations through the OPT and the LWA stations are automatically triggered to follow the VLA pointing. Specific instructions for eLWA instrument setups are provided in the OPT manual.

    The two LWA stations currently available in this mode are: LWA1 (close to the center of the VLA) and LWA-Sevilleta (on Sevilleta National Wildlife Refuge, ~80 km North-East of the VLA).

     

    Current limitations

    • VLA+LWA (eLWA) observing cannot be mixed with other observing modes in a single scheduling block/observing script. This limitation is due to the disabling of the array geometric delay model.
    • Data quality is highly susceptible to the low-frequency interference environment. Known sources of interference can be the active sun, powerline arcing, or self-generated interference from AC power components or digital electronics. The environment is regularly monitored and mitigation measures are taken, which in rare cases can significantly delay or prevent execution of observations.
    • Correlated data products are available through the LWA archive only and will eventually also be available through the regular NRAO archive. The PI of a proposal will be notified when correlated products are available.
    • Data calibration is recommended to be performed using AIPS.

    WIDAR Correlator

    Introduction

    The correlator configurations offered for general observing may be divided into three basic modes: wideband, spectral line, and subarrays. The possible setups are also subject to the integration time and data rate restrictions outlined in the section on Time Resolution and Data Rates. The possibilities and restrictions are embodied in the Resource Catalog Tool for proposing (RCT-proposing) and in the Resources section of the Proposal Submission Tool (PST), which must be used to define the correlator configuration for General Observing (GO) and Shared Risk Observing (SRO) proposals.

    Additionally, phased-array configurations are possible.  These are allowed for VLBI experiments (see the section on VLBI Observations) and for phased-array pulsar observations.

    Wideband and spectral line observing modes with the WIDAR correlator are described below. For the subarray mode, we refer to the Subarrays section of the OSS, and for pulsar observing modes, we refer to the Pulsar Observing section of the OSS.

    For technical details about the WIDAR correlator, refer to References 14 in Documentation.

     

    WIDAR Correlator: Wideband Observing

    The wideband observing setups provide the widest possible bandwidth for a given observing band, with channel spacing depending on the number of polarization products as listed in the following table 4.1.1:

     

    Table 4.1.1: Wideband & Subarray Correlator Options
    (all but P and L-bands)
    Polarization products Channel spacing
    Full (RR, RL, LR, LL) 2 MHz
    Dual (RR and LL) 1 MHz
    Single (RR or LL) 0.5 MHz

     

    8-bit wideband setups are available for all observing bands, providing a total of 2 GHz of bandwidth per polarization (1 GHz per polarization at L-band, and 256 MHz per polarization at P-band). 3-bit setups are available for all bands above S-band, providing total bandwidths per polarization of 4 GHz (C/X-bands), 6 GHz (Ku-band), or 8 GHz (K/Ka/Q-bands). In all cases, except for P and L-band, each of the subbands is 128 MHz wide. At L-band the default is 64 MHz/subband, yielding channels twice as narrow as those listed in the table above, while at P-band the default is 16 MHz/subband, resulting in 125 kHz channel spacing.

    In many frequency bands, the total processed bandwidth is less than that delivered by the front-end. In those cases, the observer may independently tune two 1 GHz baseband pairs when using the 8-bit samplers, or four 2 GHz baseband pairs when using the 3-bit samplers, or choose to have a mix 8-bit and 3-bit samplers. The tuning restrictions are described in the section on VLA Frequency Bands and Tunability, and the 8-bit and 3-bit samplers are described in the section on VLA Samplers.

     

    WIDAR Correlator: Spectral Line Observing



    Basebands and Subbands

    Observers have access to very flexible correlator configurations using up to 64 subbands in up to 4 basebands sampled with the 8-bit and/or the 3-bit samplers. These capabilities may be summarized as follows:

    • Two 1 GHz baseband pairs using the 8-bit samplers, or four 2 GHz baseband pairs using the 3-bit samplers, independently tunable within the limits outlined in the section on VLA Frequency Bands and Tunability. The 8-bit baseband pairs are referred to as A0/C0 and B0/D0, while the 3-bit samplers are A1/C1, A2/C2, B1/D1, and B2/D2. The AC/BD nomenclature corresponds to that of the IF pairs in the pre-expansion VLA.
    • Up to 16 subband pairs (spectral windows) in each 3-bit baseband pair, and up to 32 subbands in each 8-bit baseband pair, for a total of up to 64 subbands in any correlator configuration:
      • Tuning, bandwidth, number of polarization products, and number of channels can be selected independently for each subband;
      • All subbands must share the same integration time;
      • No part of a subband can cross a 128 MHz boundary;
      • Subband bandwidths can be 128, 64, 32, …, 0.03125 MHz (128 / 2n, n=0, 1, …, 12).
    • The sum over subbands of channels times polarization products is limited to 16384 (without recirculation):
      • These may be spread flexibly over subbands and polarization products, in multiples of 64: 64, 128, 192, 256, 384, …, 16384 cross-correlation products;
      • Recirculation may be used to increase the number of channels per subband for subbands narrower than 128 MHz. Baseline Board stacking may be used to increase the number of channels per subband for setups requiring less than 64 subbands;
      • Assigning many channels to a given subband may reduce the total bandwidth and/or the total number of subbands available.

    The remainder of this section discusses the various limitations in more detail, including some examples to show how they come up in practice.

     

    Subband Tuning Restrictions

    Each subband may be placed anywhere within a baseband, with the caveat that no subband may cross a 128 MHz boundary. Mathematically:

    νBB0 + n×128 MHz <= νsbLow <= νsbHigh <= νBB0 + (n+1)×128 MHz

    where:

    νBB0 the lower frequency edge of the baseband;
    n= 0, 1, …, 7 (, …, 15) (any integer between 0 and 7 for 8-bit, between 0 and 15 for 3-bit);
    νsbLow
    the lower edge of the subband
    (the subband center frequency minus half the subband bandwidth);
    νsbHigh
    the upper edge of the subband
    (the subband center frequency plus half the subband bandwidth).

    For example, if the baseband were tuned to cover 10000–11024 MHz, one could place a 64 MHz subband to cover 10570–10634 MHz, but not to cover 10600–10664 MHz because that would cross the 128 MHz boundary at 10640 MHz. Note in particular that the center of a baseband is a boundary and no line should be observed at the baseband center.

    The figure below illustrates these restrictions:

    Correlator configuration figure: bandpass8jul12.png

    The black curve shows the analog filter response for an 8-bit baseband covering 1024 MHz, starting at νBB0. The dashed blue vertical lines show the 128 MHz boundaries; no subband can cross those boundaries and 128 MHz subbands are thus constrained to cover a region between two of those boundaries, with no finer tuning being possible. Narrower subbands, like the 64 MHz subband shown here in red, can be shifted around arbitrarily within one of the 128 MHz slots, but cannot cross any of these boundaries. (The dotted vertical red lines show the boundaries of the 64 MHz subband, while the solid curve shows an illustrative line within the subband.)

    The analog filter shape defining the baseband rolls off severely at one edge of the baseband, so the 128 MHz slot at that edge has reduced sensitivity. The baseband edge is at the lowest sky frequency in the baseband when using the upper sideband, and at the highest sky frequency in the baseband when using the lower sideband.

     

    Subband Bandwidths and the Digital Filter Response

    The bandwidth for each subband may be selected independently, and can be any of 128/2n MHz, for n= 0, 1, …, 12: 128, 64, 32, 16, 8, 4, 2, or 1 MHz, or 500, 250, 125, 62.5, or 31.25 kHz.

    The usable portion of the subband is set by three effects. First, as discussed above, the analog filters which define the baseband are not perfect, leading to lower sensitivity in the 128 MHz near the baseband edge for the 8-bit samplers.

    Second, because the digital filters are not infinitely sharp, the rejected sideband leaks in at both edges of the subband. This leads to additional (aliased) noise, with a factor ~2 increase in the noise at the subband edges, dropping to a few percent within a few percent of the subband edge. The precise filter shape and noise increase is a complex but predictable function of the subband bandwidth (sbBW) and the subband tuning.

    The third effect stems from the offset frequencies used for sideband rejection in the WIDAR correlator. The local oscillators at the individual antennas are tuned to slightly different frequencies, with those offsets taken out in the correlator. This means that each antenna observes a slightly different sky frequency, and thus some baselines will not give an interesting correlation near one edge of the subband. The maximum frequency shift is currently set to 32×f0, with the fundamental f0 being set to f0 = max(25.6 kHz×sbBW/128 MHz, 100 Hz). Here sbBW is the smallest subband bandwidth within the baseband. For the wider subband bandwidths the maximum frequency shift corresponds to <1% of that bandwidth, but for narrower subbands the effect can be severe. For instance, a 31.25 kHz subband has f0 = 100 Hz, and a maximum frequency shift of 3.2 kHz—10% of the subband may be lost on some baselines.

     

    Spectral Channels and Polarization Products

    Each subband, without recirculation enabled, can have a different number of channels and polarization products, subject to two limitations:

    1. For the ithsubband, the number of spectral channels can be:
      • 64 nBlBP,i with full polarization products (RR,RL,LR,LL)
      • 128 nBlBP,i with dual polarization products (RR and LL)
      • 256 nBlBP,i with a single polarization product (RR or LL)
      Here nBlBP,i= 1, 2, 3, 4, 5, …, 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.
    2. The sum over all subbands of nBlBP,i must be less than or equal to 64, the number of Baseline Board pairs in the correlator. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64 × 256 = 16,384 (without recirculation).

    Baseline Boards are the boards in the WIDAR correlator where the actual cross-multiplications are done. There are 128 Baseline Boards arranged as 64 Baseline Board pairs (BlBPs). The limitations given here correspond to the capabilities of the individual boards and the finite number of boards the correlator has. Use of more than one pair per subband (i.e., nBlBP>1) is known as Baseline Board stacking; see additional details about this below.

    Limitation #1 corresponds to table 4.2.1 of the options for subband bandwidth and spectral resolution when using nBlBP Baseline Board pairs for a subband:

    Table 4.2.1: Subband Bandwidth and Spectral Resolution Options (without recirculation)
    Subband bandwidth &
    total velocity coverage
    Full polarization products
    (RR, RL, LR, LL)
    64nBlBP spectral channels

    Channel spacing:
    Dual polarization products
    (RR and LL)
    128nBlBP spectral channels
    Channel spacing:
    Single polarization product
    (RR or LL)
    256nBlBP spectral channels

    Channel spacing:
    128 MHz 38400/νGHz km/s 2000/nBlBP kHz 600/nBlBPGHz km/s 1000/nBlBP kHz 300/nBlBPGHz km/s 500/nBlBP kHz 150/nBlBPGHz km/s
    64 19200 1000 / nBlBP 300 / nBlBP 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP
    32 9600 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP
    16 4800 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP
    8 2400 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP
    4 1200 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 /n BlBP
    2 600 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP
    1 300 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP
    0.5 150 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP
    0.25 75 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP
    0.125 37.5 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP
    0.0625 18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
    0.0325 9.375 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP 0.122 / nBlBP 0.037 / nBlBP

    Subband bandwidth and spectral resolution options. Note that the table entries refer to the spacing between spectral channels—that spacing is before any frequency smoothing, so these channels are not independent.

    • nBlBP is the number of Baseline Board Pairs assigned to the subband.
    • Each subband may have a different number of spectral channels and polarization products, and each may be tuned independently.
    • There can be at most 16 subbands per baseband, and nBlBP must be an integer: 1, 2, 3, 4, 5, …, 64.
    • The sum of nBlBP over all subbands must be less than or equal to 64.
    • Use of more than one BlBP for a subband may further restrict the number of subbands available in one or more of the basebands; see text for details.

    Here are four examples of allowed general observing setups which use all 64 BlBPs to produce the maximum number of channels times polarization products:

    Table 4.2.2: Example BlBP Setups
    Baseband Subband

    Pol'n
    Products

    Spectral
    channels

    nBlBP
    Example 1 A0/C0 sb0 RR 16384 64
    Example 2 A0/C0 sb0 RR 8192 32
    A0/C0 sb1 RR, LL 1024 8
    A0/C0 sb2 RR, LL 512 4
    B0/D0 sb0 RR, LL 2048 16
    B0/D0 sb1 RR,RL,LR,LL 256 4
    Example 3 A0/C0 sb0 RR 8192 32
    A0/C0 sb1 LL 1024 4
    A0/C0 sb2 RR, LL 1024 8
    A0/C0 sb3 RR,RL,LR,LL 1024 16
    A0/C0 sb4 RR,RL,LR,LL 256 4
    Example 4 A0/C0 sb0-5 RR,RL,LR,LL 64 6 x 1
    A0/C0 sb6 RR, LL 3840 1 x 30
    A0/C0 sb7 RR 768 1 x 3
    A0/C0 sb8 RR,RL,LR,LL 192 1 x 3
    B0/D0 sb0-2 RR,RL,LR,LL 64 3 x 1
    B0/D0 sb3 LL 768 1 x 3
    B0/D0 sb4 RR, LL 2048 1 x 16

     

    Recirculation

    When the subband bandwidth is less than the maximum 128 MHz, the spare clock cycles that become available in the correlator hardware can be re-purposed to compute additional lags using a single baseline board. This increases the spectral resolution within the subband, and is known as recirculation. Each factor of two reduction in subband bandwidth results in an additional factor of two maximum lags; therefore for subbands of 128 MHz / N the possible spectral resolution (in units of frequency) can be increased by a factor of N2.

    Recirculation vs. Baseline Board Stacking

    When faced with the choice between recirculation and Baseline Board stacking to increase the number of channels in a subband, we recommend recirculation for subbands narrower than 128 MHz; this is supported in observatory software (RCT-proposing, OPT). Using recirculation rather than stacking frees up more Baseline Boards for other uses; alternatively the experiment becomes less dependent on all Baseline Board pairs being available/working at the time of observation. For subbands of 128 MHz, recirculation is not possible, and Baseline Board stacking must be utilized to increase the number of channels.

    The present implementation of recirculation is that, for each halving of the subband bandwidth, the number of channels in the subband may be doubled without having to use additional correlator hardware. The maximum recirculation factor for a subband is 128/(subband bandwidth in MHz) and, of course, subject to other configuration restrictions such as data rate.

    Baseline Board Stacking

    As opposed to recirculation, which increases the number of channels in a subband by exploiting otherwise unused CPU resources, Baseline Board stacking adds more channels to a subband by adding correlator hardware resources, i.e., using up more Baseline Board pairs. Using Baseline Board stacking may therefore limit the number of subbands available in one or more of the basebands. Understanding how this works requires understanding some of the details of the correlator hardware. That understanding is built into the RCT-proposing, and observers may simply use that tool to find out whether their particular setup will, in fact, work. But the results can be confusing without some understanding of the hardware constraints from which they arise. These hardware constraints are complex, and most observers will not need to understand these details. The following section is for those who are attempting complex line experiments and who find the RCT-proposing restricting the number of subbands and/or channels they can use in unexpected ways. Most observers can skip the following section.

    Baseline Board Stacking and Correlator Use

    First let us consider how the correlator hardware is organized. The cross-multiplications in the WIDAR correlator are spread across 64 Baseline Board pairs (BlBP), arranged into 4 quadrants of 16 BlBP each. Each baseband is connected directly to one of those quadrants. In the simplest mode, each of the 16 BlBP of a quadrant handles the correlations for one of the 16 subbands of the corresponding baseband. Four basebands and four quadrants are required to handle the full 8 GHz of bandwidth per polarization provided by the 3-bit (wideband) samplers: 8 GHz is split into four basebands of 2 GHz each, with each baseband fed into a different BlBP quadrant. Each BlBP in that quadrant handles a subband of maximum bandwidth 128 MHz, so 16 BlBP handles 16 subbands for a total of 16×128 MHz = 2048 MHz.

    A single BlBP produces 256 cross-correlations per baseline for a single subband, which can be used for a single polarization product (e.g., RR or LL with 256 spectral channels), or two (RR and LL with 128 spectral channels each), or four (RR, RL, LR, and LL with 64 spectral channels each).

     

    When using the 8-bit samplers, the total bandwidth is only 2 GHz per polarization, split into two basebands of 1 GHz each. The simplest continuum setup uses only two quadrants, since there are only two basebands; and only 8 subbands are required to span the 8×128 MHz = 1024 MHz of each baseband. Three-quarters of the correlator BlBP hardware remain unused.

     

    The spectral line mode allows access to these extra correlator resources through Baseline Board stacking: using multiple BlBPs to process the same subband and produce more cross-correlations for that subband. This is done using crossbar switches which make the data for a single subband available to several BlBPs. Those BlBPs can then be used to produce more spectral channels for that subband, with n BlBPs producing 256×n cross-correlations per baseline. The limit on the total number of cross-correlations (16384) stems from the total number of BlBPs (64): 64×256 = 16384.

    Unfortunately, completely flexible crossbar switches are expensive and could not be implemented in the VLA's correlator. This means that one cannot route a given subband to a randomly-chosen BlBP. The routings which are possible, are as follows:

    1. A subband in a baseband can be routed to any BlBP within the corresponding quadrant.
    2. Data coming into a given BlBP in one quadrant, can be routed to the corresponding BlBP in any other quadrant.

    Routing option #1 means that one could use all the BlBPs within a quadrant to correlate a single subband, yielding 16×256 = 4096 cross-correlations for that subband:

     

    Routing option #2 means that one could use the BlBPs in all 4 quadrants to correlate a single subband. One simple case would use 4 BlBPs to correlate each of the 16 subbands in a single baseband, yielding 4×256 = 1024 cross-correlations for each of those subbands. Note that in this case, no BlBPs are left to correlate any data from the second baseband.

    Using routing option #2 does come with a subtle cost: assigning a BlBP in quadrant X to correlate a subband corresponding to quadrant Y removes that BlBP from use in the baseband corresponding to quadrant X…and therefore also removes the corresponding subband in that baseband. So, getting more channels for a subband in one baseband may prevent the use of a subband in a different baseband. To take a simple example, consider an experiment where one wishes to observe a single line in dual polarization with 512 channels (requiring 4 BlBPs), plus as much continuum bandwidth as possible. Naively, one would say there are 16 subbands in each baseband; one is used for the spectral line, so that leaves 16+15 = 31 subbands, and with the widest subband bandwidth (128 MHz) the total available continuum should be 31×128 MHz = 3968 MHz per polarization. Actually, however, there are only 15+15 subbands available, or 30×128 MHz = 3840 MHz per polarization, because the spectral line subband has eaten one BlBP corresponding to the other baseband:

    If the same spectral line required twice as many channels, this will result in the loss of two subbands in both of the basebands:

    In some cases one may want to use a different routing to use up subbands in one baseband in preference to another. For instance, the same spectral line setup (2048 cross-correlations for a single spectral line subband, plus as much continuum as possible) could be set up to allow 13 continuum subbands in the A0/C0 baseband, and the full 16 continuum subbands in B0/D0:

    Understanding these confusing constraints can help observers set up the VLA more effectively to achieve their scientific goals. For instance, in a mixed line+continuum experiment, it works best to use the resource tools to set up the baseband tunings and subband channelization for the most important lines first, then add the desired continuum, and then see what correlator resources remain for any lines of secondary interest.

    The above examples all use BlBP pair stacking in powers of 2, but this is not required. To give some idea of more complex possibilities, the following tables (4.2.3 and 4.2.4) give two examples of other possible configurations. The RCT display shows how the Baseline Boards are used to process the individual subbands. The cyan boxes (shaded when printed out in black and white) show the Baseline Boards used to process data from baseband A0/C0, while the yellow boxes show Baseline Boards used to process data from baseband B0/D0.

    Table 4.2.3: Complex Configuration Example #1
    Baseband Subband Pol'n products Spectral channels nBlBP
    A0/C0 sb0 RR 10240 40
    A0/C0 sb1 LL 768 3
    A0/C0 sb2 RR,LL 2176 17
    B0/D0 sb0 RR 256 1
    B0/D0 sb1 RR,LL 384 3
    RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

     

    Table 4.2.4: Complex Configuration Example #2
    Baseband Subband Pol'n products Spectral channels nBlBP
    A0/C0 sb0 RR 4352 17
    A0/C0 sb1 RR, LL 1152 9
    B0/D0 sb0 RR,RL,LR,LL 192 3
    B0/D0 sb1 RR, LL 4480 35
    RCT display: corr-cfg-fig:sro2_8bit_ac17+9_bd3+35

     

    The individual subbands can have different bandwidths, and those bandwidths may be chosen completely independently of the number of spectral channels in each subband. So, for instance, a subband with a bandwidth of 2 MHz and 1152 spectral channels would have a channel separation of 2 MHz/1152 = 1.736 kHz; but the observer could equally well choose a bandwidth of 64 MHz for that subband, leading to a channel separation of 64 MHz/1152 = 55.56 kHz.

    Use of the 3-bit samplers further extends the possibilities. Here is one example:

    Table 4.2.5: 3-bit Complex Configuration Example #1
    Baseband Subband Pol'n products Spectral channels nBlBP Quadrant(s): Column(s)
    A1/C1 sb0-8 RR, LL, RL, LR 9 x 64 9 x 1 Q1: 0–8
    A1/C1 sb9 RR, LL 1 x 1152 1 x 9 Q1 & Q3: 9–11, 14 / Q4: 9
    A1/C1 sb10 RR 1 x 1792 1 x 7 Q1 & Q3 & Q4 : 12,13 / Q2: 13
    A1/C1 sb11 RR, LL 1 x 384 1 x 3 Q1 & Q2 & Q3: 15
    A2/C2 sb0-11 RR, LL, RL, LR 12 x 64 12 x 1 Q2: 0–11
    A2/C2 sb12 LL 1 x 768 1 x 3 Q2: 12, 14 / Q4: 14
    B1/D1 sb0-3 RR, LL, RL, LR 4 x 64 4 x 1 Q3: 0–3
    B1/D1 sb4 RR, LL, RL, LR 1 x 320 1 x 5 Q3: 4–8
    B2/D2 sb0-6 RR, LL, RL, LR 7 x 64 7 x 1 Q4: 0–6
    B2/D2 sb7 RR, LL 1 x 640 1 x 5 Q4: 7, 8, 10, 11, 15
    RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

     

    Once again, the RCT-proposing implements all of these constraints.

    Documentation

    Documentation

    Documentation for VLA data reduction, image making, observing preparation, etc., can be found in various manuals. Current manuals are available on-line. Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the VLA user, not listed here, are available from our website.

    1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: "Synthesis Imaging in Radio Astronomy." This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. The lectures of the 2014 workshop are available at the 14th Synthesis Imaging Workshop web site.
    2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series.
    3. GUIDE TO OBSERVING WITH THE VLA: Describes details of how to observe with the VLA once you have been allocated time on the VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide). Including special observing modes such as:
      1. CALIBRATION (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/calibration)
      2. OBSERVING WITH THE 8-BIT (up to 2 GHz bandwidth) & 3-BIT (up to 8 GHz bandwidth) SAMPLER SYSTEMS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up);
      3. SPECTRAL LINE OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/line); 
      4. HIGH FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/hifreq);
      5. LOW FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/lofreq);
      6. VERY LOW FREQUENCY OBSERVING (< 500 MHz) (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/vlofreq);
      7. POLARIMETRY (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/pol);
      8. MOSAICKING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/mosaicking);
      9. RADIO FREQUENCY INTERFERENCE (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/rfi);
      10. MOVING OBJECTS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/moving);
      11. VLBI AT THE VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/vlbi).
    4. *CASA COOKBOOK (deprecated with last updates for CASA 4.7.2): The CASA Cookbook for use of the package for data reduction of VLA (& ALMA) data is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See (http://casa.nrao.edu/docs/cookbook/)
    5. CASA Online Documentation: https://casadocs.readthedocs.io/en/stable/
    6. VLA CASA Guides: Tutorials and data reduction examples of VLA data in CASA (https://casaguides.nrao.edu/index.php/Karl_G._Jansky_VLA_Tutorials)
    7. *AIPS COOKBOOK: The Astronomical Image Processing System (AIPS) software is able to fully calibrate VLA data and do most imaging operations. The exception is the wide-band (bandwidth synthesis) deconvolution which is being developed in CASA only. ALMA data may also be reduced in AIPS although the package is not fully qualified to calibrate data from the ALMA linearly-polarized feeds. The Cookbook description for calibration and imaging under the AIPS system can be found near all public workstations in the SOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See (http://www.aips.nrao.edu/cook.html)
    8. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See (http://www.aips.nrao.edu/goaips.html)
    9. *VLA CALIBRATOR LIST: This page contains the list of VLA Calibrators in both 1950 and J2000 epoch. See (https://science.nrao.edu/facilities/vla/observing/callist)
    10. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983.
    11. *HISTORICAL VLA MEMO SERIES: archive memo series from the early days of the VLA. See (http://library.nrao.edu/vlam.shtml)
    12. *RECENT VLA MEMO SERIES: the memo series relating to the expanded capabilities of the VLA. See (http://library.nrao.edu/evla.shtml)
    13. *The VLA Expansion Project: Construction Project Book. The Expanded VLA Project Books contains the technical details of the VLA Expansion construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml.
    14. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains a useful source of information on much of the VLA. There are a few hard copies at the VLA and in the DSOC. Much of this document is now available for download (https://science.nrao.edu/facilities/vla/obsolete/green-book). Note: it does not include any information about the hardware and software specific to the expanded Karl G. Jansky VLA.
    15. WIDAR: The DRAO design and development documents of the WIDAR correlator of the VLA are available at http://www.aoc.nrao.edu/widar/docs/.

     

    Online Tools & Important Links

    The NRAO User Portal. (https://my.nrao.edu) This is a gateway to the NRAO interactive services that include the Proposal Submission Tool (PST).

    The NRAO Proposal Submission Tool (PST) online manual. (https://science.nrao.edu/facilities/vla/docs/manuals/proposal-guide/pst)

    The VLA Exposure Calculator Tool (ECT) online manual. (https://science.nrao.edu/facilities/vla/docs/manuals/propvla/determining)

    The VLA Exposure Calculator Tool (ECT). (https://obs.vla.nrao.edu/ect/)

    The Resource Catalog Tool for proposers. (https://rctp.vla.nrao.edu/rct/)

     

    Acknowledgements

    Many thanks to all the VLA staff and our RSRO participants who have worked long and hard to commission these capabilities and who have helped to create this extensively updated set of documentation.

    NRAO is grateful to Professor Rob Ivison for supporting the upgrade of some of the 3-bit samplers on the VLA via a grant from the European Research Council. For observations using the 3-bit samplers between May 2015 and March 2018 we encourage users to include the following text in the Acknowledgments section of their publications:

    "We acknowledge funding towards the 3-bit samplers used in this work from ERC Advanced Grant 321302, COSMICISM."

    Contact Information

    Please go to the People page for information on key personnel at NRAO-Socorro.

    Please direct queries to the NRAO Helpdesk; you can expect a response within one to two business days.  


     

    Editor's Notes

    This Observational Status Summary for the Karl G. Jansky (expanded) VLA is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. For questions on the content, or suggestions that would enhance the clarity of this guide, we recommend contacting the NRAO Helpdesk.

    VLA Observational Status Summary 2025A Complete Manual

    Introduction

    Purpose of Document, Older Versions of the OSS

    This Observational Status Summary (OSS) summarizes the instrumental status of the Karl G. Jansky Very Large Array (VLA) for the A-configuration, for the observing period 21 February 2025 through 25 August 2025 (Semester 2025A), and should be used when preparing proposals for the 31 January 2024 deadline. Multi-configuration proposals that include this configuration may also be submitted. Additionally, proposals requesting only configurations that will fall in future semesters may be submitted if the Principal Investigator is a graduate student. NRAO offers this service to provide scientific and technical feedback for students, and to provide them with an opportunity to re-submit their proposals for their principal semester with this information in hand.

    The OSS is intended as a ready reference for those contemplating use of the VLA for their astronomical research. The information is in summary form; those requiring greater detail should use the NRAO Helpdesk, or refer to the manuals and documentation listed in Documentation. Most of the information contained here, and much more, is available through the VLA science web page and the companion VLBA science web page. For capabilities offered in previous semesters, we refer to our overview of all OSS versions available online.

    The VLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for its efficient use. Although the NRAO strives to make using the VLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first–time users are encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. For more details, refer to the Visiting the DSOC and VLA page. The VLA is an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it.

     

    Overview of the VLA

    The Karl G. Jansky Very Large Array (VLA) is a 27–element interferometric array, arranged along the arms of an upside-down Y, which produces images of the radio sky at a wide range of frequencies and resolutions. The VLA is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico.

    The basic data produced by the VLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array's elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing, commonly called phased array, allows operation of the array as a single element through coherent summation of the individual antenna signals. This mode is most commonly used for Very Long Baseline Interferometry (VLBI) observing and for observations of rapidly varying objects, such as pulsars.

    The VLA can vary its resolution over a range exceeding a factor of ∼50 through movement of its component antennas. There are four basic antenna arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A, respectively. For details about antenna positions in the various configurations we refer to the stations position file (pdf version).

    The VLA completes one cycle through all four configurations in an approximately 16 month period. Consult the Configuration Plans and Proposal Deadlines page or recent NRAO and AAS newsletters for current and up-to-date configuration schedules and associated proposal deadlines. Refer to the Guide to Proposing for the VLA for information on how to submit an observing proposal.

    Observing projects on the VLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours with only one or, perhaps, a few target sources. However, since the VLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See the section on Snapshots for more detail.

    All VLA antennas are outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers cover the frequency ranges of 1–2 GHz (L-band), 2–4 GHz (S-band), 4–8 GHz (C-band), 8–12 GHz (X-band), 12–18 GHz (Ku-band), 18–26.5 GHz (K-band), 26.5–40 GHz (Ka-band), and 40–50 GHz (Q-band). Additionally, all antennas of the VLA have receivers for lower frequencies, enabling observations at P-band (200–500 MHz). These low frequency receivers also work at 4-band (54–86 MHz), and new feeds have been deployed on all VLA antennas to observe at this frequency range.

    The VLA correlator is both powerful and flexible. Details of the correlator configurations being offered for VLA science are described in the WIDAR Section of this document. It is important to realize that the VLA correlator is fundamentally a spectral line correlator and that even continuum observations are done in a wide-band mode with many channels.

    Offered VLA Capabilities during the Next Semester

    The Call for Proposals

    The most recent Call for Proposals summarizes the General Observing (GO) capabilities being offered for the Karl G. Jansky Very Large Array (VLA).

    In addition to these general capabilities, NRAO continues to offer shared risk observing options for those who would like to push the capabilities of the VLA beyond those offered for general use. These are the Shared Risk Observing (SRO) and Resident Shared Risk Observing (RSRO) programs.

    Details about what is being offered for each program are given below. If you have any questions or problems with any link or tool, please submit a ticket through the NRAO Helpdesk.

    Considering the lack of hybrid configurations after semester 2016A, guidelines on how to substitute such configurations with the use of principal array configurations are presented in the Array Configurations section of the Guide to Proposing for the VLA.

     

    General Observing (GO) and Shared-Risk Observing (SRO)

    Summary of Capabilities

    As described in the Call for Proposals, the VLA offers continuous frequency coverage from 1–50 GHz in the following observing bands: 1–2 GHz (L-band); 2–4 GHz (S-band); 4–8 GHz (C-band); 8–12 GHz (X-band); 12–18 GHz (Ku-band); 18–26.5 GHz (K-band); 26.5–40 GHz (Ka-band); and 40–50 GHz (Q-band). Both single pointing and mosaics with discrete, multiple field centers will be supported under General Observing (GO). In addition to these, all VLA antennas are equipped with 224–480 MHz (P-band) and 54–86 MHz (4-band) receivers near the prime focus. Data rates of up to 60 MB/s (216 GB/hour) will be available to all users as GO, combined with correlator integration time limits per band and per configuration, as described in the Time Resolution and Data Rates section. Limitations on frequency settings and tuning ranges are described in the Frequency Bands and Tunability section.

    The GO capabilities being offered are:

    Capability Description
    8-bit samplers
    • Standard full polarization default setups for:
      • 2 GHz bandwidth continuum observations at S/C/X/Ku/K/Ka/Q bands (16 × 128 MHz subbands)
      • 1 GHz bandwidth continuum observations at L-band (16 × 64 MHz subbands)
      • 256 MHz bandwidth continuum observations at P-band (16 × 16 MHz subbands)
      • 12 MHz bandwidth Stokes I continuum observations only* at 4-band (3 x 4 MHz subbands)
      • Dual 4/P-band for Stokes I continuum observations only*
    • Flexible setups for spectroscopy using two independently tunable, 1 GHz baseband pairs, each of which can be split into up to 32 flexibly tunable subbands
    • Single, dual, and full polarization products for non-default setups

    *Note: 4-band and dual 4/P-band observations are offered for Stokes I continuum only using standard full polarization default setups. Polarization, spectral-line, or the use of non-standard setups, should be submitted as a RSRO proposal.
    3-bit samplers
    • Standard full polarization default setups for:
      • 8 GHz bandwidth continuum observations at K/Ka/Q-bands
      • 6 GHz bandwidth at Ku-band
      • 4 GHz bandwidth at C/X-bands
    • Flexible setups for spectroscopy using four independently tunable, 2 GHz baseband pairs, each of which can be split into up to 16 flexibly tunable subbands
    • Single, dual, and full polarization products for non-default setups
    Mixed 3-bit and 8-bit samplers
    • Allows more flexibility for simultaneous continuum and high-resolution spectral line observing

    Subarrays

    • Up to 3 independent subbarrays using standard 3-bit continuum setups, or a mix of standard 3-bit and standard 8-bit continuum setups, and up to 3 independent subarrays with changing standard continuum setups in a given subarray (e.g., to perform reference pointing at X-band for high frequency observations).

    Y27 or Y1 for VLBI

    Solar observing

    • All solar observing except the L-band reverse-coupled system*

    On-The-Fly Mosaicking (OTF)

    • P-, L-, S-, and C-bands only, using linear interpolations in Equatorial Coordinates; no subarrays

    Pulsar

    • Phase-binned and coherent-dedispersion (YUPPI) pulsar observing, except 4-band YUPPI and combining YUPPI mode with VLBI recording, which are Resident Shared Risk Observing (RSRO) capabilities.

    *Note: The VLA L-band (1-2 GHz) has a special signal path (the "reverse coupler" path) that allows coherent radio bursts to be observed without saturating the system, as the brightest of these solar bursts can exceed 105 solar flux units, or 109 Jy. This signal path has not yet been fully commissioned and is therefore not yet available under GO.

    SRO capabilities can be set up via the Observing Preparation Tool (OPT) and run through the dynamic scheduler without intervention, but are not as well tested as GO capabilities. Data rates higher than 60 MB/s (216 GB/hour) and up to 100 MB/s (360 GB/hour) are considered SRO. A summary of the SRO capabilities being offered are:

    • On-the-Fly (OTF) mosaicking for X-, Ku-, K-, Ka-, and Q-bands (used when each pointing on the sky is on the order of several seconds or less), but not using subarrays.
    • Wideband VLA for VLBI: Enables recording of VLA WIDAR continuum-mode correlations during VLA phased array (Y27) VLBI observations. Currently, this only supports standard VLA 8-bit continuum modes with a 2-GHz bandwidth. See the VLBA Call for Proposals for more details.
    • eLWA: Joint LWA and VLA 4-band observations using a single 8 MHz subband centered at 76 MHz, and 4-bit VDIF output. Note: During semester 2025A, the LWA is expected to be undergoing infrastructure upgrades and availability of the telescopes (LWA1 and LWA-SV) may be limited.  Those interested in using this mode should contact Greg Taylor at gbtaylor@unm.edu for more details.

    We expect that most SRO programs will have no or only minor problems that can be corrected quickly. If an SRO program fails, however, and it becomes clear that detailed testing with additional expertise is needed, then the project must make an experienced member from their team available to help troubleshoot the problem. In some cases, this may require the presence of that experienced member in Socorro. If adequate support from the project is not given, then the time on the telescope will be forfeited. The additional effort is to be determined based on discussions with the NRAO staff and management and the project team.

    The guidelines for General and Shared Risk observing proposals, along with information about tools and other advice, can be found in the VLA Proposal Submission Guidelines.

    Resident Shared Risk Observing (RSRO)

    Summary of Capabilities

    The VLBA Resident Shared Risk Observing (RSRO) program provides users with early access to new capabilities in exchange for a period of residency in Socorro to help commission those capabilities.

    RSRO proposals should be submitted using the NRAO Proposal Submission Tool in response to a regular proposal call. The proposal should include a scientific justification, as for normal proposals, which will be peer reviewed as part of NRAO's time allocation process. Selecting "VLA RSRO" from the "Observing Mode" menu on the Resources page makes an "RSRO Comments" text-entry facility available for describing the technical resources required. A description of the personnel who will be involved in the effort along with their expertise and availability should also be included in the technical justification.

    We emphasize the "shared risk" nature of the RSRO program. Since observers will be attempting to use capabilities under development and in the process of being commissioned, NRAO can make no guarantee of the success of any observations made under this program, and no additional commitment is made beyond granting the hours actually assigned by the peer review process.

    Proposals for any area of user interest bit offered under GO or SRO are welcome. Here, we provide some examples of capabilities that are being utilized in recent RSRO proposals.

    • Correlator dump times shorter than 50 msec, including integration times as short as 5 msec for transient detection, or data rates above 100 MB/s. In order to reduce the data rate, frequency averaging in the correlator may be utilized in RSRO proposals;
    • YUPPI pulsar mode combined with VLBI recording;
    • Subarray observations with setups other than the default continuum setups, or observations with more than 3 subarrays;
    • Currently, OTF observing is implemented as linear interpolations in Equatorial Coordinates (i.e., RA/Dec). This can be expanded to allow using stripes linear in Galactic coordinates (l,b), as well as more complex patterns other than linear in RA and Dec, such as Rosetta or Spiral patterns. Note that these must still adhere to the restrictions of the OTF mode under General Observing, i.e., using the full array below 8 GHz (up to C-band), and no subarrays.

    The guidelines for Resident Shared Risk Observing proposing, along with requirements and considerations, can be found in the VLA Proposal Submission Guidelines.

     

    Commensal Observing Systems at the VLA

    There are three commensal systems on the VLA that may take data at the same time as your proposed observation. The first is the VLITE system, which will take data at P-band during regular observations that use bands other than P-band. Hence, VLITE is turned off by default during P-band or dual 4/P-band observations. The VLITE system is deployed on up to eighteen VLA antennas. Observers wishing to gain access to the commensal VLITE data taken during their VLA observations should follow the instructions on the VLITE web page for doing so. The second is the realfast system, which takes data at very fast dump rates in an effort to detect Fast Radio Bursts (FRBs). This system is fully commissioned for observing at L- through X-bands, in parallel with standard continuum correlator configurations. The third commensal system, COSMIC SETI, enables the search for extraterrestrial intelligence (SETI) using the VLA, and collects data during unconflicted PI science observations. For information about commensal observing see the Commensal Observing with NRAO Telescopes page.

    To report errors or problems encountered in any link or while using any NRAO tool listed here, please submit a ticket through the NRAO Helpdesk.

    Performance of the VLA during the Next Semester

    Resolution

    Resolution

    The VLA's resolution is generally diffraction-limited, and thus is set by the array configuration and the observing frequency. Like all synthesis arrays, the VLA is sensitive only to structures on a range of angular scales between the diffraction limit (the smallest angular scale detectable) and a "Largest Angular Scale" (which depends on the fringe spacing formed by the shortest baselines in the configuration). For emission structures smaller than the diffraction limit (θ ∼ λ/Bmax), the VLA acts like a single-dish instrument—the resulting image is smoothed to the resolution of the array. For emission structures larger than the detectable range, the VLA is simply blind to the emission; this is a limitation unique to interferometers. No subsequent processing can fully recover the missing information from these large scales. It can only be obtained by observing in a more compact VLA array configuration or with data from an instrument that is sensitive to the missing angular scales, such as a large single dish or a compact array of smaller antennas.

    Table 3.1.1 displays the VLA's resolution and the scale at which severe attenuation of large-scale structure occurs. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power; the resolution element) for the central frequency for each band, and the largest angular scale of detectable emission.

    Table 3.1.1: Configuration Properties
    ConfigurationABCD
    Bmax (km1) 36.4 11.1 3.4 1.03
    Bmin (km1) 0.68 0.21 0.0355 0.035
    Band Synthesized Beamwidth θHPBW(arcsec)1,2,3
    74 MHz (4) 24 80 260 850
    350 MHz (P) 5.6 18.5 60 200
    1.5 GHz (L) 1.3 4.3 14 46
    3.0 GHz (S) 0.65 2.1 7.0 23
    6.0 GHz (C) 0.33 1.0 3.5 12
    10 GHz (X) 0.20 0.60 2.1 7.2
    15 GHz (Ku) 0.13 0.42 1.4 4.6
    22 GHz (K) 0.089 0.28 0.95 3.1
    33 GHz (Ka) 0.059 0.19 0.63 2.1
    45 GHz (Q) 0.043 0.14 0.47 1.5
    Band Largest Angular Scale θLAS(arcsec)1,4
    74 MHz (4) 800 2200 20000 20000
    350 MHz (P) 155 515 4150 4150
    1.5 GHz (L) 36 120 970 970
    3.0 GHz (S) 18 58 490 490
    6.0 GHz (C) 8.9 29 240 240
    10 GHz (X) 5.3 17 145 145
    15 GHz (Ku) 3.6 12 97 97
    22 GHz (K) 2.4 7.9 66 66
    33 GHz (Ka) 1.6 5.3 44 44
    45 GHz (Q) 1.2 3.9 32 32
    These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith.
    Notes:
    1. Bmax is the maximum antenna separation, Bmin is the minimum antenna separation, θHPBW is the synthesized beam width (FWHM), and θLAS is the largest angular scale structure visible to the array.
    2. The listed resolutions are appropriate for sources with declinations between −15 and +75 degrees.
    3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θHPBW. The values for snapshots are about 1.3 times the listed values.
    4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations, the quoted numbers should be divided by two.
    5. For the C configuration, an antenna from the middle of the north arm is moved to the central pad N1. This results in improved imaging for extended objects, but may slightly degrade snapshot performance. Note that although the minimum spacing is the same as in D configuration, the surface brightness sensitivity and image fidelity to extended structure is considerably inferior to that of the D configuration.

    The following figure is a graphical representation of the synthesized beamwidths for natural and robust weighting for the four main array configurations between 1 and 50 GHz. Also available are synthesized beamwidth figures for the low frequency (1–12 GHz) and the high frequency (12–50 GHz) receiver bands.

    Sensitivity

    Sensitivity

    The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by:

    where:

    - SEFD is the system equivalent flux density (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEFD indicate more sensitive performance. For the VLA's 25–meter paraboloids, the SEFD is given by the equation SEFD = 5.62TsysA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band.
    - ηc is the correlator efficiency (~0.93 with the use of the 8-bit samplers).
    - npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in RCP or LCP.
    - N is the number of antennas.
    - tint is the total on-source integration time in seconds.
    - Δν is the bandwidth in Hz.

    Figure 3.2.1 shows the SEFDs as a function of frequency used in the VLA exposure calculator for those Cassegrain bands currently installed on VLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Figure 3.2.2 shows the SEFDs as a function of frequency for the P-band; these measurements are based on imaging of a field far from the galactic plane. Table 3.2.1 gives the SEFD at some fiducial VLA frequencies.

     Figure 3.2.1: SEFD used in the Exposure Calculator for the VLA. Left: The system equivalent flux density as a function of frequency for the L, S, C and X-band receivers. Right: The system equivalent flux density as a function of frequency for the Ku, K, Ka, and Q-band receivers. SEFDs at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere and were determined under good conditions. 

    Figure 3.2.2: The SEFD used in the VLA Exposure Calculator as a function of frequency for the P-band receiver

    Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical:

    • For the more commonly used robust weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values.
    • Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 3.2.1 shows the confusion noise in D configuration assuming robust weighting (see Condon et al. 2012, ApJ, 758),  which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 3.2.1; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This confusion primarily affects low frequency observations.
    • Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission: primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band.
    • Losses from the 3-bit samplers. The VLA's 3-bit samplers incur an additional 10–15% loss in sensitivity above that expected—i.e., the efficiency factor ηc = 0.78 to 0.83.

     

    Table 3.2.1: SEFDs and D-Configuration Confusion Limits
    Frequency   SEFD
    (Jy)
     

    RMS confusion level
    in D config (µJy/beam)
    Robust Weighting

    0.39 GHz (P) 2790 5330
    1.5 GHz (L) 420 74
    3.0 GHz (S) 370 12
    6.0 GHz (C) 310 2
    10.0 GHz (X) 250 negligible
    15 GHz (Ku) 320 negligible
    20 GHz (K) 500 negligible
    33 GHz (Ka) 600 negligible
    45 GHz (Q) 1300 negligible

     

    In general, the zenith atmospheric opacity to microwave radiation is very low: typically less than 0.01 at L, C, and X-bands; 0.05 to 0.2 at K-band; and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Observing conditions are best at night and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable.

    Observers should remember that clouds, especially clouds with large water droplets (thunderstorms), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz.

    Tipping scans—which are currently unavailable but will be implemented at some time in the future—can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the complex gain calibrator (amplitude and phase) and target source are observed.

    When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following items:

    • Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target.

      Both CASA and AIPS allow the application of elevation-dependent gains and an estimated opacity generated from ground-based weather through the CASA tasks gencal and plotweather, and AIPS task INDXR.

    • Pointing. The SEFD quoted above assumes good pointing. Under calm, nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime can be much worse—occasionally exceeding 1 arcminte due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 11 mph (5 m/s). Wind speeds near the stow limit 45 mph (20 m/s) will have a similar negative effect at 8 and 15 GHz.

    To achieve increased pointing accuracy, referenced pointing is recommended where a nearby calibrator is observed in interferometric pointing mode every hour or so. The local pointing corrections measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2–3 arcseconds (but more typically 5–7 arcseconds) if the reference source is within about 15 degrees in azimuth and elevation of the target source and the source elevation is less than 70 degrees. At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup.

    Use of referenced pointing is highly recommended for all Ku, K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.3 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution.

    To aid VLA proposers there is an online guide to the exposure calculator; the exposure calculator provides a graphical user interface to these equations.

    Special caveats apply for P-band (230–470 MHz) observing. The SEFD's in Figure 3.2.2 or that listed in table 3.1.2 are from an observation taken far from the Galactic plane, where the sky brightness is about 30K. At P-band, Galactic synchrotron emission is very bright in directions near the Galactic plane. The system temperature increase due to Galactic emission will degrade sensitivity by factors of two to three for observations in the plane, and by a factor of five or more at or near the Galactic center. Additionally, the antenna efficiency (currently about 0.31 for 300 MHz) will decline with both increasing and decreasing frequencies from the center of P-band.

    The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by:

    where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼1.5λ/Bmax), and S in mJy per beam, the parameter F depends on the synthesized beam, therefore on the array configuration, and has the approximate value of F = 190, 18, 1.7, 0.16 for A, B, C, and D configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ΔIm, for S. Note that Equation 2 is a beam-averaged surface brightness; if a source size can be measured, then the source size and integrated flux density should be used in Equation 2 and the appropriate value of F calculated. In general, the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 6 of Reference 1, listed in Documentation.

    For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791):

    where D is the distance to the galaxy in Mpc, and SΔV is the HI line area in units of Jy km/s.

     

    VLA Frequency Bands and Tunability

    Bands

    For observations taken with the 8-bit samplers, each receiver can tune to two different frequencies, each 1024 MHz wide, within the same frequency band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies, except for the low-band receiver (50–500 MHz), which provides linear polarization (X and Y). Each of these four data streams follows the VLA nomenclature and are known as IF (for Intermediate Frequency channel) A, B, C, and D. IFs A and B provide RCP (or Y when applicable), IFs C and D provide LCP (or X when applicable). IFs A and C are always at the same frequency, as are IFs B and D (but note that the A and C IFs frequency is usually different from the B and D frequency). We normally refer to these two independent data streams as IF pairs, i.e., the A/C pair and the B/D pair. In 8-bit mode, a maximum of 1024 MHz can be correlated for each IF pair (see the WIDAR Section), for a total maximum bandwidth of 2048 MHz. To distinguish this 8-bit system from the 3-bit system, these IF pairs are denoted A0/C0 and B0/D0.

    More options are available with the 3-bit samplers. This system provides four (R,L) polarization pairs, each 2048 MHz wide. The A/C IF pair provides two sampled pairs, labelled A1/C1 and A2/C2, and the B/D IF pair provides two sampled pairs, labelled B1/D1 and B2/D2.

    For more details on the 8-bit and 3-bit samplers see the VLA Samplers section.

    The tuning ranges, along with default frequencies for continuum applications, are given in Table 3.3.1 below.

    Table 3.3.1: Default frequencies for continuum applications
    Band Range1 8-bit continuum applications (GHz) 3-bit continuum applications (GHz)
      (GHz) IF pair A0/C0 IF pair B0/D0 IF pair A1/C1 IF pair A2/C2 IF pair B1/D1 IF pair B2/D2
    4 m (4) 0.054 – 0.0842 .070 – .082
    90 cm (P) 0.20 – 0.503 0.224 – 0.4803
    4 m (4)/90 cm (P) 0.054 – 0.0842/
    0.20 – 0.503
    0.224 – 0.4803 .070 – .082
    20 cm (L) 1.0 – 2.04 1.0 – 1.54 1.5 – 2.04
    13 cm (S) 2.0 – 4.0 2.0 – 3.0 3.0 – 4.0
    6 cm (C) 4.0 – 8.0 4.5 – 5.5 5.5 – 6.5 4.0 – 6.0 6.0 – 8.0
    3 cm (X) 8.0 – 12.0 8.0 – 9.0 9.0 – 10.0 8.0 – 10.0 10.0 – 12.0
    2 cm (Ku) 12.0 – 18.0 13.0 – 14.0 14.0 – 15.0 12.0 – 14.0 14.0 – 16.0 16.0 – 18.0
    1.3 cm (K) 18.0 – 26.5 20.2 – 21.2 21.2 – 22.2 22.0 – 24.0 24.0 – 26.0 18.0 – 20.0 20.0 – 22.0
    1 cm (Ka) 26.5 – 40.0 32.0 – 33.0 31.0 – 32.0 33.0 – 35.0 35.0 – 37.0 29.0 – 31.0 31.0 – 33.0
    0.7 cm (Q) 40.0 – 50.0 40.0 – 41.0 41.0 – 42.0 44.0 – 46.0 46.0 – 48.0 40.0 – 42.0 42.0 – 44.0

    Notes:

    1.  Listed here are the nominal band edges. For all bands, the receivers can be tuned to frequencies outside this range, but at the cost of diminished performance. Contact the NRAO Helpdesk for further information.
    2. The 4-band system is currently under development. The default frequency range maximizes sensitivity of the system and provides a nominal bandwidth of 12 MHz and a channel resolution of 64 kHz.
    3. The default setup for P-band will provide 16 subbands from the A0/C0 IF pair, each 16 MHz wide, to cover the frequency range 224–480 MHz. The channel resolution is 125 kHz. 
    4. The default frequency setup for L-band comprises two 512 MHz IF pairs (each comprising 8 contiguous subbands of 64 MHz) to cover the entire 1–2 GHz of the L-band receiver.

     

    Tuning Restrictions

    In general, for all frequency bands except Ka, if the total span of the two independent IF pairs of the 8-bit system (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For K, Ka, and Q-bands—the only bands where a span greater than 8 GHz is possible—there are special rules:

    • At Ka-band, the low frequency edge of the A0/C0 IF pair must be greater than 32.0 GHz. There is no restriction on the B0/D0 frequency, unless the B0/D0 band overlaps the A0/C0 band when the latter is tuned at or near the 32.0 GHz limit. In this case, the Observation Preparation Tool (OPT) may not allow the requested frequency setups. Users wanting to use such a frequency setup are encouraged to contact the NRAO Helpdesk for possible tuning options.
    • At K and Q-bands, if the frequency span is greater than 8.0 GHz, the B0/D0 frequency must be lower than the A0/C0 frequency.

    For the 3-bit system, the maximum frequency span permitted for the A1/C1 and A2/C2 IF pairs is about 5000 MHz. The same restriction applies to B1/D1 and B2/D2. The tuning restrictions given above for the separation and location of the 8-bit pairs A0/C0 and B0/D0 also apply to the 3-bit pairs, with A0/C0 replaced by A1/C1 and A2/C2, and B0/D0 replaced by B1/D1 and B2/D2.

     

    VLA Samplers

    The VLA is equipped with two different types of samplers, 8-bit with 1GHz bandwidth, and 3-bit with 2GHz bandwidth. The choice depends on your science goals and on technicalities described below.

    The 8-bit Set consists of four 8-bit samplers running at 2.048 GSamp/sec. The four samplers are arranged in two pairs, each pair providing 1024 MHz bandwidth in both polarizations. The two pairs are denoted A0/C0 and B0/D0. Taken together, the four samplers offer a maximum of 2048 MHz coverage with full polarization. The frequency spans sampled by the two pairs need not be adjacent. Some restrictions apply, depending on band, as described in the section on Frequency Bands and Tunability.

    The 3-bit Set consists of eight 3-bit samplers running at 4.096 GSamp/sec. The eight samplers are arranged as four pairs, each pair providing 2048 MHz bandwidth in both polarizations. Two of these pairs, denoted A1/C1 and A2/C2 cannot span more than 5000 MHz (lower edge of one to the higher edge of the other). The same limitation applies to the second pair, denoted B1/D1 and B2/D2. The tuning restrictions are described in the section on Frequency Bands and Tunability. Taken together, the eight 3-bit samplers offer a maximum of 8192 MHz coverage with full polarization.

     

    Which set to use?

    • S, L, and 4/P-band observations, whether line or continuum, should use the 8-bit sampler set.
    • C and X-band continuum observations should use 3-bit samplers in order to exploit the full 4 GHz bandwidth: in spite of the 15% reduction in sensitivity that comes with 3-bit (at equal bandwidth to the 8-bit samplers—see below for details) and the reduced effective bandwidth after removing RFI, this still provides superior overall sensitivity. For more details we refer to EVLA memo 166.
        • Note: C-band is impacted by strong RFI caused by microwave links near 6 GHz in the A and B configurations. As a result, 3-bit data obtained with the standard setup are corrupted. We advise observers to use mixed 3-bit and 8-bit samplers. For more details, refer to the VLA Observing Guide.
    • Ku, K, Ka, and Q-band continuum observations should use the 3-bit samplers for maximum bandwidth.
    • Wide-band spectral line searches requiring more than 2 GHz span should use the 3-bit samplers.
    • Spectral-line observations which fit within two, possibly disjoint, 1 GHz bands should use the 8-bit set.
    • Simultaneous continuum and high resolution spectral line observation can use mixed 3-bit and 8-bit samplers. The 3-bit samplers in this case will be set up to deliver the continuum data, while the 8-bit samplers will be for the spectral line data. This mix mode can be used in C-band and higher.


    Major Characteristics of each Set

    The 8-bit samplers are warranted for observations at 4/P, L, and S-bands. The full analog bandwidth from the receivers fits within the 2048 MHz span covered by the samplers.

    For the 3-bit samplers, users need to be aware of the following issues:

    • Sensitivity: compared to the 8-bit system, the sensitivity of the 3-bit samplers is worse by ~15% (at equal bandwidth). Alternatively, a given continuum noise level requiring on-source integration time T with the 8-bit (two bands of 1GHz), requires 0.33T with the 3-bit (4 bands of 2GHz, assuming the bandwidth is available from the front end).
    • Resonances: each of the eight 3-bit samplers on an antenna has a resonance about 3 MHz wide. Each resonance is independent of all others, so there is no correlated signal between antennas. The resonance degrades the spectrum in its narrow frequency range, but has little effect on continuum observing. Bandpass solutions will be affected, but can be interpolated over. Spectral-line calibration and images at the affected frequencies will show significant loss in sensitivity. The resonances are easily seen in autocorrelation spectra, and it is recommended that users, especially spectral-line users, utilize these to locate the compromised frequencies.
    • Amplitude Calibration: The traditional method for both 8- and 3-bit systems is to observe a flux-density calibrator, use self-cal to determine the antenna amplitude calibration factors (gains), and transfer the gains to the phase calibrator and target. For 3-bit samplers this procedure gives results good to 5% between elevations of 20–70 degrees. (Expect worse at the upper edge of Q-band and/or during bad weather.) The switched power data can be used to correct for system gain variations and works well for the 8-bit samplers. For 3-bit samplers, the Pdif depends on the Psum, i.e., Pdif is non-linear and its application will bias the resulting visibilities by 5–10%. The origin of this effect is understood, but we have not yet determined how best to compensate for it. Because of this, we do not recommend use of the Psum and Pdif data to calibrate visibilities from the 3-bit samplers. We do, however, recommend that the requantizer gains in the switched power data be applied to remove gain changes. For more information about the switched power, Psum, and Pdif, see EVLA memo 145.

     

    Setting up the 8-bit or 3-bit Samplers

    Either set requires an initial scan for each individual LO (frequency) tuning, during which power levels are optimized.

    For the 8-bit system, a dummy scan of 1 minute duration is sufficient for each tuning. This  is usually done while the antennas are slewing at the start of an observing file, as the pointing direction of the antennas is not critical.

    For the 3-bit system, the requirements are more demanding, see the section on 3-bit setup within the Guide to Observing with the VLA. The minimum setup time is 1 minute for each tuning to adjust the power levels and bandpass slopes across the 2GHz samplers. These values are retained and applied if the tuning is re-encountered in the same observation. Additionally, every time the LO setup is changed—whether or not it is new (e.g., changing from 8-bit X-band reference pointing back to target)—a scan of 30 seconds is needed to reset the subband gains (requantizers) in the correlator. For better amplitude calibration at high frequencies, the 3-bit initial setup should be near the elevation of the target, so do it after the first 8-bit setup described above. For 3-bit observing without 8-bit (e.g., C or X-band without reference pointing), the power variation with elevation is small, so the 3-bit setup can be done at any elevation.

    For settings that use a mix of 3-bit and 8-bit samplers, the guidelines to set up the 3-bit samplers should be followed.


    Other issues

    The overhead for setup of 3-bit samplers can eat into observing time, especially for projects with many different LO settings, and/or sources all over the sky accompanied by band change, reference pointing, and requantizer reset for each direction. The impact is most severe for short scheduling blocks.

    Polarization testing conducted so far indicates no degradation of performance by using the 3-bit samplers.

    Field of View

    Primary Beam

    The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is θPB = 42/νGHz for frequencies between 1 and 50 GHz (L- through Q-band). At P-band the approximate value is θPB = 50/νGHz. New precise measurements of the primary beam shape have been reported in EVLA Memo 195; these allow for the correction of the primary beam attenuation in wide-field images. Both AIPS and CASA (5.0 and later versions) have these new parameters incorporated.

    With the wide-bandwidths of the VLA it is necessary to account for the variation of the primary beam with frequency in order to achieve high-dynamic range images. For this and other imaging details we refer to the Limitations on Imaging Performance section of the OSS.

    To achieve good sensitivity with a single-pointing observation, observers should take care to ensure that their targeted patch of sky fits within the primary beam (θPB) corresponding to the highest frequency of their observing band. If that is not possible, multiple overlapping pointings can be used to construct images of larger regions of sky through a technique known as mosaicking. Guidelines for mosaicking with the VLA are given in the Guide to Observing with the VLA.

    Note: The Largest Angular Scale (LAS) that can be imaged by the array is independent of the Primary Beam's field of view or the use of mosaicking to increase the field of view. A table of the band- and configuration-dependent LAS is presented in the Resolution section of this document.

     

    Chromatic Aberration (Bandwidth Smearing)

    The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array.

    These effects can be parameterized by the product of the fractional bandwidth (Δν/ν0) with the source offset in synthesized beamwidths (θ0HPBW). Table 3.5.1 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter and should be used to determine how much spectral averaging can be tolerated when imaging a particular field.

    Table 3.5.1: Reduction in Peak Response Due to Bandwidth Smearing
    (Δν/ν0)*(θ0HPBW)   Peak   Width
    0.0 1.00 1.00
    0.50 0.95 1.05
    0.75 0.90 1.11
    1.0 0.80 1.25
    2.0 0.50 2.00

    Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). Δν/ν0 is the fractional bandwidth; θ0HPBW is the source offset from the phase tracking center in units of the synthesized beam.

    Note: The VLA correlator supports frequency averaging for single subarray and non-OTF observations. Currently this capability is limited to an averaging by a factor of 2 or 4  and only for wide-band continuum science projects (appropriate for C-band through Q-band observations). Observers interested in this capability should consult the EVLA memo 199 to assess the suitability of the frequency averaging in the correlator for their observations, because the extent of the bandwidth smearing is heavily dependent on the frequency averaging factor, the array configuration, and the observing frequency.

     

    Time-Averaging Loss

    The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parametrized, except for observations at δ = 90°. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with Δν/ν0 replaced by ωeΔtint, where ωe is the Earth's angular rotation rate, and Δtint is the averaging interval.

    For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 in Documentation considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 3.5.2, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θΔtint)2, where θ is the distance from the phase center and Δtint is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see the section on Time Resolution and Data Rates) in the A and B configurations.

    Table 3.5.2: Averaging Time for a Given Amplitude Loss
    Amplitude loss
    Configuration   1.0%   5.0%   10.0%
    A 2.1 4.8 6.7
    B 6.8 15.0 21.0
    C 21.0 48.0 67.0
    D 68.0 150.0 210.0

    Note: The averaging time (in seconds) results in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines.

     

    Note: For both the chromatic aberration and the time-averaging loss, the issue is not a simple reduction in amplitude for sources far from the phase center, but a convolution to the extent that a point source far from the phase center will become resolved due to bandwidth and/or time smearing. Furthermore, the description given above for the bandwidth smearing is based on the assumption that the radiation is monochromatic to parameterize the smearing, and does not take into account the consequences of having wide-bandwidths as is the case for the VLA. Therefore, while proposing and planning for VLA observations, and depending on the objectives of the science and the location of the sources of interest within the field, including confusing sources which may be far outside the science field, the above noted guidelines need to be used to conservatively estimate the proper channel width and the correlator integration time in order to minimize the effects of the bandwidth smearing and the time smearing, respectively.

     

    Non-Coplanar Baselines

    The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the VLA with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at 90 cm and 20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations.

    Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS task IMAGR and CASA task clean. Refer to the package help files for these tasks, or consult with the NRAO Helpdesk for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the W-projection method available in CASA (see EVLA Memo 67 for more details).

    Time Resolution and Data Rates

    The default integration times for the various array configurations and frequency bands are as follows:

    Table 3.6.1: Default Integration Times
    Configurations Observing
    Bands
    Default
    integration time
    A, B, C, D 4 P 2 seconds
    A L S C X Ku K Ka Q 2 seconds
    B L S C X Ku K Ka Q 3 seconds
    C, D X Ku K Ka Q 3 seconds
    C, D L S C 5 seconds

    Observations with the 3-bit (wideband) samplers, when applicable, should use these integration times. Observations with the 8-bit samplers may use shorter integration times, but these must be requested and justified explicitly in the proposal, and obey the following restrictions:

    Table 3.6.2: Minimum integration times and maximum data rates
    Proposal type

    Minimum integration time

    Maximum data rate
    General Observing (GO) 50 msec up to 60 MB/s (216 GB/hr)
    Shared Risk Observing (SRO) 50 msec > 60 MB/s (216 GB/hour) and up to 100 MB/s (360 GB/hour)
    Resident Shared Risk Observing (RSRO) < 50 msec > 100 MB/s (360 GB/hr)

    Note that integration times as short as 5 msec and data rates as high as 300 MB/s can be supported for some observing, though any such observing is considered Resident Shared Risk Observing. For these short integration times and high data rates there will be limits on bandwidth and/or number of antennas involved in the observation. Those desiring to utilize such short integration times and high data rates should consult with NRAO staff.

    The maximum recommended integration time for any VLA observing is 10 seconds.

    Observers should bear in mind the data rate of the VLA when planning their observations. For Nant antennas and integration time Δt, the data rate is:

    Data rate ~ 45 MB/sec × (Nchpol/16384) × Nant × (Nant − 1)/(27×26) / (Δt/1 sec)
    ~ 160 GB/hr × (Nchpol/16384) x Nant × (Nant − 1)/(27×26) / (Δt/1 sec)
    ~ 3.7 TB/day × (Nchpol/16384) × Nant × (Nant − 1)/(27×26) / (Δt/1 sec)

    Here Nchpol is the sum over all subbands of spectral channels times polarization products:

    Nchpol = Σi Nchan,i × Npolprod,i

    where Nchan,i is the number of spectral channels in subband i, and Npolprod,i is the number of polarization products for subband i (1 for single polarization [RR or LL], 2 for dual polarization [RR and LL], 4 for full polarization products [RR, RL, LR, LL]). This formula, combined with the maximum data rates given above, imply that observations using the maximum number of channels currently available (16384) will be limited to minimum integration times of ~2 seconds for standard observations, and 0.8 seconds for shared risk observations.

    We note that frequency averaging in the correlator will reduce the total number of channels. Therefore, the data rate and the data volume will be reduced by the same channel averaging factor. See the Chromatic Aberration section for more details on the frequency averaging in the correlator and to assess its impact on your science.

    These data rates are challenging for transfer and analysis. Data may either be downloaded via ftp over the Internet, or shipped on hard drives for large data sets or for those with slow Internet connections (please review the data shipping policy). For users whose science permits, the Archive Access Tool allows some level of frequency averaging in order to decrease data set sizes before ftp; note that the full spectral resolution will be retained in the NRAO archive for all observations.

    Note: The data rate formula given above does not account for the auto-correlations delivered by WIDAR. Precise data rate values can be obtained through the use of the Resource Catalog Tool for proposing (RCT-proposing).

    Radio Frequency Interference

    The very wide bandwidths of the VLA mean that radio frequency interference (RFI) will be present in a far larger fraction of current observations than in observations made with the old systems. Considerable effort has gone into making the VLA's electronics as linear as possible, so that the effects of any RFI will remain limited to the actual frequencies at which the RFI exists. Non-linear effects, such as receiver saturation, should occur only for those very unlikely, and usually very brief, times when the emitter is within the antenna primary beam.

    RFI is primarily a problem within the low frequency (C, S, L, and the low-band system) bands, and is most serious in the D configuration. With increasing frequency and increasing resolution comes an increasing fringe rate, which is often very effective in reducing interference to tolerable levels.

    The bands within the tuning range of the VLA which are protected for radio astronomy are: 73.0-74.6 MHz, 322.0-328.6 MHz, 1400–1427 MHz, 1660.6–1670.0 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No significant external interference should occur within these bands.

    VLA staff periodically observes the entire radio spectrum with the VLA, from 1.0 through 50.0 GHz with 125 kHz channel resolution, to monitor the ever-changing RFI spectrum. Users concerned about the precise frequencies of strong RFI, and the likelihood of being affected, are encouraged to peruse these plots. To access these plots, or for more information on RFI, including the impact of satellite transmissions, please see the RFI section in the Guide to Observing with the VLA.

    Subarrays

    The continuum subarray option offers two 1 GHz baseband pairs with the 8-bit samplers in up to 3 subarrays or four 2 GHz baseband pairs with the 3-bit samplers, with the same spectral channel and polarization product options as are available for wideband observing. The setup for each subarray is completely independent in terms of observing frequency, polarization products, and integration times.

    When using three subarrays, there are some restrictions on the number of antennas in each subarray. The Baseline Board in the correlator treats each set of 4 antennas independently, using a separate column of correlator chips. With 8 such columns, the correlator can handle up to 8×4 = 32 antennas. The correlator configuration software requires that a given column not be split across subarrays. This does not matter when using only two subarrays, but forces some subtle restrictions when using three. For instance, one cannot observe with 9 antennas in each of 3 subarrays, because 9 antennas requires three columns (two with 4 antennas each, and one with 1 antenna); three subarrays of 9 antennas each would require 3×3 = 9 columns, one more than are actually available. Splitting the array into 10, 9, and 8 antennas is allowed, since the first two subarrays use 3 columns each, while the third uses only two.

    Table 3.8.1 gives four examples of how correlator resources can be split into multiple subarrays. Antennas in each subarray are color-coded: red for subarray 1, green for subarray 2 (if present), and blue for subarray 3 (if present). The last column gives the number of antennas in each subarray (e.g., in the setup shown in the first row, subarray 1 has 10 antennas, subarray 2 has 9 antennas, and subarray 3 has 8 antennas). In all cases a total of 27 antennas are used. The columns are numbered in reverse order (C7 to C0) to match the numbering scheme used on the actual Baseline Boards.

    Table 3.8.1: Some Possible Subarray Options
    Number of antennas correlated using each Baseline Board column Number
    of antennas
    C7 C6 C5 C4 C3 C2 C1 C0
    4 4 2 4 4 1 4 4 10 + 9 + 8
    4 4 4 2 4 4 4 1 14 + 13
    4 4 3 4 4 3 3 2 11 + 11 + 5
    4 4 4 4 4 4 3 27

    For more information on subarrays, please see the Subarray section in the Guide to Observing with the VLA.

    Positional Accuracy & Astrometry

    The position of a target can be determined to a small fraction of the synthesized beam, limited by atmospheric phase stability, the proximity of an astrometric calibrator, the calibrator-source cycle time, and the SNR on target.

    In preparation for observing, the a priori position must be known to within the antenna primary beam, except perhaps for mosaicking observations. In the special case of using the phased VLA as a VLBI element, the a priori position must be accurate to within the synthesized beam of the array.

    In post-processing, target positions are typically determined from an image made after phase calibration, i.e., correcting the antenna and atmospheric phases as determined on the reference source. The accuracy of the calibration determines the accuracy of the positions in the image. Note that phase self-calibration imposes the assumed position of the model, i.e., makes the position indeterminate. Therefore, an absolute position cannot be determined after self-calibration, but relative positions between features within a self-calibrated image are valid.

    It may help to think of astrometry as two methods, narrow-field and wide-field.

    Narrow-field astrometry

    In narrow-field astrometry, the target is close to the phase tracking center and the antennas nod every few minutes between the target and a calibrator.  If no special calibration provisions are taken, under typical conditions, an astrometric accuracy of ~10% of the synthesized beam FWHM can often be obtained.  For example, an observation in Ka-band (~33 GHz) in A-configuration might reach an astrometric accuracy of ~10 milliarcseconds (mas).  When care is taken (special calibrations and ideal observing conditions), the accuracy can approach 1–2% of the synthesized beam, with a floor of ~2 mas.  If such accuracies are needed, we strongly recommend obtaining advice from VLA staff in setting up the observations.

    Astrometric calibrators are marked J2000 A in the VLA calibrator list, and have a positional accuracy of ~2 mas. Other catalogs from the USNO and the VLBA are also useful, but offsets may exist between the VLA and VLBA centroids arising from extended structure in the particular source and the different resolutions of the arrays.

    For studies of proper motion and parallax, the absolute accuracy of a calibrator may be less important than its stability over time. Close, or in-beam calibrators with poor a priori positions, can be used and tied to the ICRF reference frame in the same or separate observations.

    Phase stability can be assessed in real time from the Atmospheric Phase Interferometer (API) at the VLA site, which uses observations of a geostationary satellite at ~12 GHz. Dynamic scheduling uses the API data to run a project under suitable conditions specified by the user. VLBI projects using the phased VLA will typically be fixed date and not dynamically scheduled.

    Wide-field astrometry

    Wide-field astrometry is used to determine the positions of targets within the primary beam, referenced to a calibrator within the beam or close by. In addition to the previous effects, there are distortions as a function of position in the field, from small errors in the Earth Orientation Parameters (EOP) used at correlation time, differential aberration, and phase gradients across the primary beam. With no special effort, the errors build up to roughly one synthesized beam at a separation of ~104 beams from the phase tracking center. Not all these errors are fully understood, and accurate recovery of positions over the full primary beam in the wide-band, wide-field case is a research area. These effects are handled somewhat differently in the post-processing packages. Check with VLA staff for more details via the NRAO Helpdesk.

    Limitations on Imaging Performance

    Imaging Fidelity

    Image fidelity is a measure of the accuracy of the reconstructed sky brightness distribution. A related metric, dynamic range, is a measure of the degree to which imaging artifacts around strong sources are suppressed, which in turn implies a higher fidelity of the on-source reconstruction.

    With conventional external calibration methods, even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is most often the effective phase stability of the telescope due to atmospheric/ionospheric fluctuations, although pointing errors and changes in atmospheric opacity may also be a limiting factor. If a good model of the sky brightness distribution exist (e.g., use of compact structures of sufficient strength—though a good model of resolved sources in the field of view may also be used), standard self-calibration can be counted on to improve the images. At low frequencies, where the dominant phase error is due to ionospheric plasma density fluctuations, more advanced techniques may be required to account for change of ionospheric phase across the field of view. Depending on the underlying nature of the errors, dynamic ranges in the thousands to hundreds of thousands can be achieved using these techniques. With the WIDAR correlator and wide-band receivers, i.e. large bandwidths resulting in high sensitivities, self-calibration methods can be extended to observations of sources with much lower flux densities than ever possible with the old VLA.

    The choice of image reconstruction algorithm also affects the correctness of the on-source brightness distribution. The CLEAN algorithm is most appropriate for predominantly point-source dominated fields. Extended structure is better reconstructed with multi-resolution and multi-scale algorithms. For high dynamic ranges with wide bandwidths, algorithms that model the sky spectrum as well as the average intensity can yield more accurate reconstructions.

     

    Invisible Structures at Large Angular Scales

    An interferometric array acts as a spatial filter so that, for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Image reconstruction algorithms such as multi-resolution and multi-scale CLEAN may help reduce these negative bowls, but care must be taken in choosing appropriate scale sizes to work with.

    Table 3.1.1 in the Resolution section gives the largest scale visible to each configuration/band combination.

     

    Poorly Sampled Fourier Plane

    Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in Documentation.

     

    Sidelobes from non-Deconvolved Sources

    At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (tclean) are well suited to this task. Sidelobe confusion is a strong function of observing band—affecting most strongly L and P-band observations—and is rarely a significant problem for observations at frequencies above 4 GHz.

     

    Sidelobes from Strong Sources

    An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral-line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere.

     

    Wide-band Imaging

    The very wide bandpasses provided by the VLA enable imaging over 2:1 bandwidth ratios: at L, S, and C-bands, the upper frequency is twice that of the lower frequency. It is this wide bandwidth which enables sub-microJy sensitivity.

    In many cases, where the observation goal is a simple detection and there are no strong sources near to the region of interest, standard imaging methods that combine the data from all frequencies into one single image (multi-frequency-synthesis) may suffice. A rough rule of thumb is that—provided a strong source is not adjacent to the target zone—if the necessary dynamic range in the image is less than 1000:1 (i.e., the strongest source in the beam is less than 1000 times higher than the noise), a simple wide-band map may suffice.  

    For higher dynamic ranges, complications arise from the fact that the brightness in the field of view dramatically changes as a function of frequency, both due to differing structures in the actual sources in the field of view and due to the attenuation of the sources by the primary beam. One symptom of such problems is the appearance of radial spokes around bright sources, visible above the noise floor, when imaged as described above.  

    The simplest solution is to simply make a number of maps (e.g., one for each subband) which can then be suitably combined after correction for the primary beam shape. But with up to 64 subbands available with the VLA's new correlator, this is not always the optimal approach, both from a practical standpoint and frequency-dependent image reconstruction uncertainties. Further, images at all bands must be smoothed to the angular resolution at the lowest frequency before any spectral information can be extracted. With a 2:1 bandwidth, the difference in angular resolution across the band will be significant.

    A better approach is to process all subbands simultaneously, utilizing software which takes into account the possibility of spatially variant spectral index and curvature and knows the instrumentally-imposed attenuation due to the primary beam. Such wideband (plus wide-field) imaging algorithms are now available within CASA as part of the tclean task.

     

    Wide-field Imaging

    Wide-field observing refers to both the non-coplanar nature of the VLA when observing in non-snapshot mode and the effect of the primary beam patterns.

    At high angular resolutions and low frequencies, standard imaging methods will produce artifacts around sources away from the phase center due to the non-coplanar nature of the VLA. Faceted imaging (AIPS, CASA) and W-projection (CASA) techniques can be used to solve this problem. Community image packages also employ similar techniques, usually geared toward addressing low frequency and wide-field imaging problems and might work with VLA data as well. 

    Wide field imaging also requires the accurate representation of primary beam patterns and their use during imaging. This is relevant only for very high dynamic ranges ( > 10000:1 ) or when there are very strong, confusing sources at and beyond the half-power point of the primary beam. This problem is worse with a wide-band instrument simply because the size of the primary beam—and the radius at which the half-power point occurs—varies with frequency while there is increased sensitivity out to a wider field of view. In CASA, the tclean task contains algorithms to deal with these effects by modeling and correcting for polarization squint, frequency-dependent and rotating primary beams, per antenna, during imaging. Work is under way to develop full-polarization primary beam corrections.  Please note, however, that most advanced methods are computationally very demanding, will lead to a significant increase in processing time, and may not always be required. Therefore, in the interest of practicality, they should be used only if there is evidence of artifacts without these methods. At low frequencies, wide-field imaging typically includes both non-coplanar effects plus primary beams whereas at high frequencies it is primarily primary beam effects.

    Mosaicking is another form of wide-field imaging in which data from multiple pointings are combined either during imaging as a joint mosaic (CASA) or after imaging by stitching several images together (CASA, AIPS). It is also possible to apply a- and w-term corrections for joint mosaicking in CASA, however in most situations this requires significant computational resources and cannot be provided through standard NRAO post-processing computing. Furthermore, the current implementation in tclean requires for this case particular attention to run efficiently in a cluster environment. For more and up-to-date information on this topic, please refer to the CASA documentation on synthesis imaging.

     

    Calibrating the Flux Density Scale

    Normal calibration of the flux density scale for VLA observations is effected by including a scan on a source of presumed known flux density in each Scheduling Block (SB). Using that known flux density source, the flux density of the complex gain calibrator(s) can be determined and then transferred to your target source(s). Historically, 3C48 and 3C286 have been the standard sources for which NRAO has assumed flux densities are known as a function of frequency, and which have been recommended as flux density scale calibrator sources for the VLA. Restrictions on baseline length as a function of VLA configuration and observing band were supplied which, if followed, allowed relatively accurate flux density scale calibration. We have recently improved the ability to calibrate the flux density scale by providing sky brightness models for these sources in CASA and AIPS, which loosens the restrictions on configurations and bands. We have also added the sources 3C138** and 3C147 to the list of calibrators that have models. However, 3C48*, 3C138**, and 3C147 have spectral flux densities that vary with time (3C286, along with 3C295 and 3C196, are constant), so some care should be taken if the most accurate flux density scale calibration is desired.

    Note: While accurate models are available in both AIPS and CASA for various frequency bands for the calibrators 3C286, 3C48*, 3C147, and 3C138**, neither 3C295 nor 3C196 has such models in CASA. Therefore, the VLA CASA calibration pipeline will fail if these two calibrators are used. Furthermore, 3C295 and 3C196 may not be suitable for all VLA configurations and frequencies even if one chooses to not use the pipeline.

    A single observation of a few minutes of one of the above-mentioned flux density scale calibrators will suffice for most observers. If possible, the flux density scale calibrator should be observed at a time when it is nearly at the same elevation as the complex gain calibrator, especially for the highest four bands (Ku-, K-, Ka-, and Q-band). This is not always possible because of timing and geometry of sources, and that it is not typically known when an SB will be executed (so elevations versus time are uncertain). Flux density scale calibration accuracy in this case should be of order 10% at 4- and P-bands, 5% at L- through Ku-bands, and 10-15% for the three higher bands. If more accuracy is needed, a more careful strategy should be adopted, potentially using multiple flux density scale calibrators. The fundamental accuracy of the scale is ~5% at 4- and P-bands, 3% at L- through Ku-bands, increasing to 5% at Q-band. See Perley and Butler (2017) for more details on how the spectral flux densities of 3C48*, 3C138**, 3C147, and 3C286 (and many other sources) have been determined across the frequency range from 50 MHz to 50 GHz and how they vary versus time, along with information on the fundamental accuracy of the flux density scale when using these sources.

    If less accuracy is needed in the flux density scale calibration, an observation of one of these standard sources need not necessarily be included in an SB. As an example, for a short triggered observation where a simple detection is desired, the time spent slewing back and forth to the flux density scale calibrator can make the SB significantly longer than it could otherwise be. In this scenario, the switched power measurement can be used to calibrate the flux density scale; see EVLA Memo 120 for some background. This technique, which should only be used with the 8-bit samplers, is not a standard path of calibration, but it is possible. The flux density scale accuracy in this case is ~10% for L- through Ku-bands, increasing to ~20% at Q-band; not nearly as good as using the "standard" method of flux density scale calibration, but it may be sufficient for some observers.

    For reference, the polynomial expression for the spectral flux density for 3C286 determined in Perley and Butler (2017) is: \[\log(S) = 1.2481 - 0.4507 \log(f) - 0.1798 \log^2(f) + 0.0357 \log^3(f)\] where S is the flux density in Jy, and f is the frequency in GHz.

    The tables below show flux densities determined using the polynomial coefficients for a few sources at a single frequency within each of the VLA bands.

    Flux densities (Jy) of Standard Calibrators for January 2016
    Source 75 MHz 350 MHz 1500 MHz 3000 MHz 6000 MHz 10000 MHz 15000 MHz 22000 MHz 33000 MHz 45000 MHz
    3C48* = J0137+3309 72.8 42.2 15.4 8.44 4.42 2.68 1.79 1.22 0.815 0.601
    3C138** = J0521+1638 26.5 16.1 8.25 5.44 3.39 2.33 1.72 1.28 0.949 0.761
    3C147 = J0542+4951 58.0 52.3 21.0 12.0 6.45 3.99 2.73 1.93 1.39 1.13
    3C196 = J0813+4813 129 44.4 13.6 6.98 3.38 1.91 1.20 0.763 0.473 0.329
    3C286 = J1331+3030 30.0 25.9 14.6 9.91 6.39 4.50 3.37 2.54 1.88 1.49
    3C295 = J1411+5212 124 58.4 21.2 11.0 5.06 2.70 1.60 0.970 0.571 0.385
    Flux densities (Jy) of Standard Calibrators for January 2019
    Source 328 MHz 1465 MHz 2565 MHz 4885 MHz 6680 MHz 11320 MHz 16564 MHz 25564 MHz 32064 MHz 48064 MHz
    3C48* = J0137+3309 43.9 15.6 9.82 5.48 4.12 2.56 1.86 1.33 1.11 0.816
    3C138** = J0521+1638 15.9 8.26 6.00 4.00 3.23 2.24 1.69 1.25 1.06 0.821
    3C147 = J0542+4951 53.9 21.4 13.8 7.88 5.91 3.67 2.61 1.82 1.53 1.14
    3C196 = J0813+4813 46.5 13.8 8.14 4.22 3.00 1.67 1.08 0.656 0.508 0.313
    3C286 = J1331+3030 25.8 14.6 10.9 7.33 5.97 4.12 3.15 2.30 1.92 1.44
    3C295 = J1411+5212 61.1 21.6 12.8 6.42 4.45 2.33 1.43 0.819 0.596 0.405

    We refer the reader to the VLA Observing Guide for the practical considerations (e.g., observing frequency and array configuration, as well as post processing) regarding the choice of the flux density scale calibrator in their scheduling blocks.

    * The flux density scale calibrator 3C48 has been undergoing a flare since January 2018 or so.  While we have not fully characterized this with the VLA, other instruments have measured it at some frequencies. At Ku-band the magnitude of the flare is of order 10%.  The effect will be smaller at lower frequencies (of order 5% at L-band), and might be larger at higher frequencies (of order 20% at Q-band).  If you care about the flux density scale of your observations at that level, you may want to re-calibrate your data once new time-variable values have been put into CASA and AIPS.

    ** The flux density scale calibrator 3C138 is currently undergoing a flare. From VLA calibration pipeline results, we have noticed that 3C138 is deviating from the model. The amount of this deviation is still being investigated by NRAO staff, but does seem to effect frequencies of 10 GHz and higher. At K and Ka-bands the magnitude of the flare is currently of order 40-50% compared to Perley-Butler 2017 flux scale. If you care about the flux density scale of your observations above 10 GHz, monitoring datasets are publicly available in the archive under project code TCAL0009, from which you may find an updated flux density ratio to use for your data.

    Complex Gain Calibration

    General Guidelines for Complex Gain Calibration

    Adequate complex gain calibration (tracking amplitude and phase fluctuations as a function of time) is a complicated function of source-calibrator separation, frequency, array scale (configuration), and weather. Since what defines adequate for some experiments is completely inadequate for others, it is difficult to define simple guidelines to ensure adequate phase calibration. However, some general statements remain valid most of the time. These are given below.

    • Under decent conditions with no thunderstorms or ionospheric storms, tropospheric effects dominate at frequencies higher than about 4 GHz; ionospheric effects dominate at frequencies lower than about 4 GHz.
    • Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the C and D configurations at L and S-bands, and in the D configuration at X and C-bands. For these cases, calibration need only be done to track instrumental changes—three or four times per hour is usually sufficient for tracking the system gains.
    • If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than a couple of times per hour at low frequencies or 15 minutes at high frequencies in order to track pointing or other effects that might influence the amplitude scale. The enhanced sensitivity of the VLA guarantees, for full-band continuum observations, that every field will have enough background sources to enable phase self-calibration at L and S-bands. At higher frequencies, the background sky is not sufficient, and only the flux of the target source itself will be available.
    • In principle, the smaller the source-calibrator angular separation, the better (even if the closer calibrator is weaker). However, the final choice will depend on the observing frequency. If deciding between a nearby calibrator with an S code in the calibrator database, and a more distant calibrator with a P code, for low frequencies (L-band and below) the nearby calibrator is usually the better choice (low frequency strategy). For higher frequencies it is advisable to use the further away but higher quality P-code calibrator (high frequency strategy). A detailed description of calibrator codes is available in the calibrator list.
    • Phase stability often deteriorates dramatically after about 10AM due to small-scale convective cells set up by solar heating, even in clear and calm conditions, especially in the summer. Observers should consider a more rapid calibration cycle for observations between this time and a couple of hours after sundown.
    • At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is needed to track tropospheric phase fluctuations if the target cannot be self-calibrated. See Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) (below).

     

    Rapid Phase Calibration and the Atmospheric Phase Interferometer

    For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-calibrator observing cycles as short as 40 seconds can be used for very small source-calibrator separations. Observing efficiency declines, however, for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. This is represented in the Observation Preparation Tool as a loop of source-calibrator scans with short scan length. This technique stops tropospheric phase variations at an effective baseline length of ∼vat/2, where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec) and t is the total switching time. Short source-calibration scans have been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Note that at a typical wind velocity in the compact D-configuration, this effective baseline length is the same as—or larger than—the longest baseline in the array and it is not worth the increased overhead of short cycle times. Under these circumstances, it is sufficient to calibrate every 5-10 minutes to track the instrumental changes. The fast switching technique will not work in bad weather (such as rain showers or when there are well-developed convection cells (thunderstorms)). It is important to correctly specify the required tropospheric phase stability as measured by the Atmospheric Phase Interferometer at observe time (see below).

    Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from References 9 and 10 in Documentation.  Also see the High Frequency Strategy guide for additional recommendations on observing at high frequencies.

    An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase. The API uses an interferometer of two, 1.5 meter antennas, separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data are heavily used for the dynamic scheduling of the VLA.

    Characteristic seasonal averages are shown in Table 3.12.1 below:

    Table 3.12.1: Seasonal API/wind values at the VLA
    Month

    API (night)
    [deg]

    API (median)
    [deg]

    API (day)
    [deg]

    Wind (night)
    [m/s]

    Wind (median)
    [m/s]

    Wind (day)
    [m/s]

    January 2.3 2.8 3.6 1.6 1.9 2.3
    February 2.9 3.4 4.5 4.0 4.3 4.5
    March 2.8 3.7 5.5 3.4 3.9 4.7
    April 3.3 4.5 6.2 5.3 5.5 5.8
    May 2.9 4.6 6.7 2.6 3.2 3.7
    June 3.8 5.5 7.4 2.5 3.9 6.3
    July 6.2 8.3 10.5 2.9 2.9 3.0
    August 5.4 7.1 11.3 1.7 2.3 3.0
    September 5.2 6.6 8.8 2.3 3.0 3.6
    October 4.2 5.3 7.4 2.3 2.9 3.7
    November 2.6 3.0 4.0 1.2 2.5 1.6
    December 2.8 3.2 4.1 1.2 1.6 2.7

    Day indicates sunrise to sunset values; night indicates sunset to sunrise values.

     

    Other Issues that Affect Complex Gain (Amplitude)

    There are other instrumental effects that cause fluctuations in gain over time for VLA observations - we describe three of them here.  Note that all of these effects are to the gain amplitude, not phase.

    Gain Curves

    The VLA antennas have elevation-dependent gain variations which are important to account for at the four highest frequency bands. Gain curves are determined by VLA staff per antenna and per band, and the necessary corrections can be applied to the visibility data in both AIPS and CASA. Additionally, atmospheric opacity will also cause an elevation-dependent gain which is also particularly notable at these four highest frequency bands. We currently do not have an atmospheric opacity monitoring procedure; users should utilize the appropriate tasks available in both AIPS and CASA to estimate and correct for the opacity using ground-based weather data. If your complex gain calibrator is near your target source, and the flux density scale calibrator is also observed at a similar elevation, then most of this elevation-based gain will be calibrated correctly during normal calibration. Note also that a good procedure for removing elevation-based gain dependencies uses the AIPS task ELINT. This task will generate a 2nd order polynomial gain correction utilizing your own calibrator observations. This will remove both the antenna and opacity gain variations, and has the decided advantage of not utilizing opacity models or possibly incorrect antenna gain curves. Use of this procedure is only practical if your observations span a wide range in elevation.

    Antenna Pointing Offsets

    Another important gain variation effect at the four highest frequency bands is that due to antenna pointing offsets. Daytime observations on sunny days can suffer pointing errors, primarily in elevation, of up to one arcminute. This effect can be largely removed by utilizing the referenced pointing procedure which determines the pointing offset of a nearby calibrator. This offset is then applied to subsequent target source observations. It is recommended that this local offset be determined at least hourly, more often during sunrise or sunset, utilizing an object within 15 degrees of the target source—preferentially at an earlier right ascension. Studies show that the maximum pointing error will be reduced to about 7 arcseconds or better if proper referenced pointing is utilized.

    Electronic Gain Variations

    The VLA's post-amplifiers are not temperature stabilized and exhibit significant gain (amplitude) changes between night and day, particularly at the four highest frequency bands. Changes as large as 30% have been seen between night and day in calm, clear conditions. These gain changes, and others caused by possible changes in attenuator settings, are monitored and can be removed by application of the internal calibration signal, whose results are recorded in the switched power table in both AIPS and CASA. These corrections are not applied by the calibration pipeline—users who wish to correct for these gain changes must utilize the appropriate tasks in AIPS or CASA. For the most accurate flux density bootstrapping, this table must be applied to the visibility data before calibration.

      Polarization

      For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle or a single observation of an unpolarized source. The complex gain calibrator chosen for the observations can also double as a polarization calibrator, provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15–50 degrees during a several hour track).

      The minimum condition that will enable accurate polarization calibration from a polarized source, in particular with unknown polarization, is three observations of a bright source spanning at least 60 degrees in parallactic angle (schedule four scans, if possible, in case one is lost); if at all possible, it is strongly recommended that five or more observations covering 100 degrees (or more) of parallactic angle in roughly uniform steps be run. If a bright, unpolarized, unresolved source is available, and known to have very low polarization, then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for small objects as compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission; 3C48 can be used at frequencies of ~3 GHz and higher, or 3C147 at frequencies over ~10 GHz. Note that 3C48 and 3C138 are variable—the polarization properties are known to be changing significantly over time, most notably at the higher frequencies (for details see Perley and Butler (2013b)).

      More information on polarization calibration strategy can be found in the Polarimetry section in the Guide to Observing with the VLA.

       

      VLBI Observations

      The VLA can participate in VLBI observations with the VLBA. This is possible by either using the VLA as a phased array (Y27) or using one of its dishes (single dish: Y1). We note that currently P-band cannot be phased. For more details see the VLBI at the VLA documentation. In phased array mode, the program TelCal derives the antenna-based delay and phase corrections needed for antenna phasing in real time. This correction is applied to the antenna signals before they are summed, requantized to 2-bits, and recorded in VDIF format on the Mark5C disk at the VLA site. The disk(s) are then transported to Socorro, NM and correlated on the DiFX correlator with other VLBI stations which participated in the observation.

      Standard VLA data, i.e., correlations between VLA antennas, are also archived in the NRAO science data archive. By default, a maximum of 512 MHz dual polarization is phased/recorded and sent to the archive. This leaves a large portion of the possible VLA bandwidth unused (up to 2 or 8 GHz total depending on sampler choice). Increasing the bandwidth to the maximum that the VLA can provide has obvious benefits for most continuum observations. Pulsar gating can be performed commensally with the phased VLA, avoiding the need for extra single dish observations. Line VLBI observations that run through the WIDAR correlator currently produce each requested VLBI subband bandwidth divided into full polarization products with 64 frequency channels, regardless of the requested spectral resolution obtained in the VLBI correlation. By adding unused baseline boards to the VLBI-specified line subbands, a closer match in spectral resolution can be offered for the standalone VLA data. If you think your science could benefit from these capabilities, please review our documentation on VLBI at the VLA or contact the helpdesk for further guidance.

      Snapshots

      The two-dimensional geometry of the VLA allows a snapshot mode whereby short observations can be used to image relatively bright, unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest.

      Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm and longer wavelengths in the D configuration. For instance, at 20 cm, a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. The deconvolution of the data is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing Faint Images of the Radio Sky at Twenty-centimeters survey (FIRST, http://sundog.stsci.edu/top.html) (B configuration) or the NRAO VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) (D configuration, all-sky).

      Shadowing and Cross-Talk

      Observations at low elevation in the C and D configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling when using the default inputs with CASA tasks importasdm and importevla. AIPS task UVFLG can be used to flag VLA data based on shadowing, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other subarrays. The CASA task flagdata can also be used to flag data based on shadowing. For more information on shadowing, please see the Antenna Shadowing section in the Guide to Observing with the VLA.

      Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. This effect is important at low frequencies in compact configurations. Careful examination of the visibilities is necessary to identify and remove this form of interference. The affected data would show time-variable high-amplitude points.

      Combining Configurations and Mosaicking

      Any single VLA configuration will allow accurate imaging of a range of spatial scales determined by the shortest and longest baselines. For extended and structured objects, it may be required to obtain observations in multiple array configurations. It is advisable that the frequencies used be the same for all configurations to be combined. The ideal combination of arrays results in a uv-plane with all cells equally filled by uv-points. To first order, this can be achieved by using the beam sizes of the individual arrays to inversely scale the on-source integration time. This approach is equivalent to achieving the same surface brightness sensitivity for all arrays on all scales. For the VLA, observations in the different configurations generate beam sizes that decrease by factors of three, i.e., C configuration generates a three times smaller beam than D configuration, B three times smaller than C, and A three times smaller than B. Thus, on-source integrations would increase by about an order of magnitude between each array. Such a drastic increase is very expensive and, in fact, not necessary since some spatial scales are common to more than a single array, which is equivalent to some uv-cells being filled more than others. The best way to fill the uv-plane depends on many factors such as declination of the source, LST time of the observation, and bandwidth.

      Experience shows for the VLA that a factor of about three in on-source integration time for the different array configurations works well for most experiments. For example, a 20min on-source time in D, 1hr in C, 3hrs in B, and 9hrs in A should produce a decent map. Using large bandwidths and multi-frequency synthesis will broaden all uv tracks radially and one may need even fewer array configurations or shorter integration times between the different arrays.

      Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicking. The VLA has no limit on the number of pointings for each mosaic. Typically hexagonal, rectangular, or individual pointing patterns are used and the overlap regions will result in an improved rms over each individual pointing. Given the many, potentially short observations, it is important to obey the data rate limits outlined in the Time Resolution and Data Rates Section. In addition to discrete or pointed mosaics, on-the-fly (OTF) mosaics (i.e. dumping the data while moving the telescopes across the source) are also available.

      Time-variable structures, such as the nuclei of radio galaxies and quasars, cause special, but manageable, problems. See the article by Mark Holdaway in Reference 2 of the Documentation for more information.

      Guidelines for mosaicking with the VLA are given in the Guide to Observing with the VLA.

      Pulsar Observing

      The VLA can be used for several kinds of pulsar observing: phase-binning using the WIDAR correlator, using the phased-array for single-beam pulsar processing in either search or fold modes, or simply standard imaging mode with fast integrations. Both phase-binning and phased-array (YUPPI) modes are available under General Observing (GO). The only exception is the 4-band YUPPI which is a Resident Shared Risk Observing (RSRO) capability. For any questions not addressed here regarding the capabilities of these observing modes, please contact the NRAO Helpdesk.

      Phased-array pulsar processing

      The "Y" Ultimate Pulsar Processing Instrument (YUPPI) is a software suite that runs in the correlator backend (CBE) computer cluster and can process a single-beam phased/summed-array data stream for pulsar observations in real time, into either folded profiles or search mode (filterbank) output. Coherent dedispersion can be optionally applied in either mode.

      In the phased-array pulsar processing mode, the voltage data streams from each antenna are divided into a number of frequency subbands within the correlator, then summed and requantized before being output to the cluster for pulsar processing. The limitation on bandwidth comes primarily from the available network connections between the correlator and cluster. In all cases, a maximum of 64 subbands total can be processed. Depending on the number of bits chosen, this results in the following total bandwidth constraints:

      Table 3.18.1: Pulsar Observing Bandwidth Constraints

      Subband bandwidth

      Subband quantization

      Max total bandwidth

      Samplers

      32 MHz 8 bits 2048 MHz 8-bit
      64 MHz 4 bits 4096 MHz 3-bit
      128 MHz 2 bits 8192 MHz 3-bit
      <32 MHz 8 bits 64*BWsub 8-bit

      As described in the VLA Frequency Bands and Tunability section, the 8-bit samplers provide two independently tunable 1 GHz IFs, while the 3-bit samplers provide four tunable 2 GHz IFs. 

      The pulsar-specific processing is done in real time using the DSPSR software package and, in principle, any processing option supported by DSPSR can be used; this will be constrained by the real-time computing power available in the cluster. In general, each subband can be divided into an arbitrary (2n) number of channels; 1 (summed), 2 or 4 detected polarization products can be output; and coherent dedispersion can be enabled or not.

      Fold mode

      In fold mode, the data are averaged modulo a known pulsar ephemeris (provided via a standard TEMPO/TEMPO2 "par file") into pulse profiles. The data can also be folded at a constant topocentric period, for example at 10 Hz to detect the injected noise cal signal. Fold integration times as short as 1 second have been tested. Up to 16384 profile bins can be used. The data are recorded in PSRFITS format using the standard 16-bit data encoding. This means the final output data rate is given by:

      Data rate = 2 bytes × Nsubband × Nchannel × Nbin × Npoln / Tint

      If the desired data rate exceeds ~25 MB/s, additional testing ahead of time may be required.

      Search mode

      In search mode, the data are simply detected and averaged over a specified amount of time before being output to disk, resulting in a filterbank data array (power vs time and frequency). Coherent dedispersion at a known DM can optionally be enabled for this.  Data can be recorded using 2, 4, 8, 16 or 32 bits, resulting in a final data rate of:

      Data rate = (Nbit/8) bytes × Nsubband × Nchannel × Npoln / Tint

      The maximum sustained output rate in this mode should be kept less than ~400 MB/s.

      Subarrays

      It is possible to use any of the phased-array pulsar modes listed here in a subarray observation, following the guidelines described in the Subarrays section. In addition, the above constraints on the pulsar processing apply to the total of all simultaneously-used subarrays, rather than each subarray separately. For example, the total number of subbands in use across all subarrays must not be greater than 64; the total (not per-subarray) data rate must meet the above constraints, etc. It is possible to use different parameters such as subband bandwidth, number of bits, or processing mode (fold versus seach) in the different subarrays.

      VLBI

      It is possible to use phased-array pulsar processing as part of a VLBI experiment; see the VLBI Observations section, and links therein, for additional information about VLBI at the VLA. The main constraint on this type of observation is that a subband that is being recorded for VLBI can not be sent to the pulsar processing system. However, since VLBI recording typically only requires a small number of subbands (2 through 8), any additional subbands produced by WIDAR can be sent to the pulsar system, following the constraints above. This provides a high time resolution data stream covering wider bandwidth than the VLBI data. One typical use case is to detect a pulsar using the VLA data stream and determine a short-term timing ephemeris covering the observation. This can then be used to gate the VLBI correlation, reducing uncertainties associated with extrapolating existing timing solutions or obtaining time on other telescopes for these purposes.

      Gated or binned visibilities

      The WIDAR correlator has the capability to internally integrate (fold) visibilities into 1 or more pulse phase bins, following a standard TEMPO-compatible pulsar ephemeris. This mode can be used to image the emission from a pulsar of known period anywhere in the telescope field of view. This provides both higher signal-to-noise ratio on the pulsar than a standard image, and allows the pulsed emission to be separated from continuous emission from other sources in the field.

      The following constraints apply to binning-mode observations:

      • Binning is limited to the case where the pulse period is divided evenly into bins covering the full pulse period.  "Gating" style observations (common in VLBI) where a single on-pulse bin is used are not supported.
      • The maximum number of pulse phase bins is 1000.
      • There is a tradeoff between the total bandwidth and the minimum bin width (pulse period divided by number of bins):
        • With 4 subbands (up to 512 MHz total), the minimum allowed bin width is 12.5 μs.
        • With 16 subbands (up to 2048 MHz total), the minimum allowed bin width is 50 μs.
        • With 64 subbands (up to 8192 MHz total), the minimum allowed bin width is 200 μs.
      • The number of channels per subband is currently limited to 128 maximum.  Combining recirculation and binning is not allowed.
      • Integration (dump) time must be an integer number of pulse periods.
      • The data rate produced in this mode is the standard VLA data rate (see the Data Rate section) multiplied by the number of bins.  The data rate must be kept less than 60 MB/s.

      It should also be noted that there is currently very limited support for binned observations in standard data processing software (e.g., CASA). Development of data analysis procedures is ongoing and users of this mode should be aware that this will likely involve some advanced/low-level manipulation of raw VLA data sets.

      Fast-dump visibilities

      While not specifically a pulsar mode, standard visibility data can be dumped as fast as 5 ms, which may be sufficient for imaging of slow pulsars. See the Time Resolution and Data Rates section for more details.

      Solar Observing

      Observations of the Sun can currently be performed in five frequency bands: 1-2 GHz (L band), 2-4 GHz (S band), 4-8 GHz (C band), 8-12 GHz (X band), and 12-18 GHz (Ku band). To observe the Sun, any 8-bit resource using these bands can be used in Solar mode (see below). As the Sun is moving across the sky, the source position should be defined in the source catalog.

      If the center of the solar disk needs to be tracked, select the Sun using the drop-down for "Solar System Body with Internal Ephemeris". Otherwise, to allow for specifying the location of interest on or near the solar disk and the differential rotation model to be used for tracking the source, solar observations require selection of a source position type of "Solar System Body with Uploaded Ephemeris". An ephemeris file can be generated using the guidelines given in the Moving Object section or, for example, by using the ALMA Solar Ephemeris Generator

      The user must select the “Solar” scan mode when specifying scan details under Observation Preparation in the OPT. This scan mode ensures that the necessary attenuators are switched into the signal path when observing the Sun and that special noise cal sources are employed in the switched power system; this allows users to flux calibrate their data. Two observing modes are currently supported:

      • For quiet Sun observations, when no strong active regions or flares expected, users should select a scan mode of “Solar Attenuators with Low Noise Internal Cal”.
      • If the science objective involves observing a strong active region and/or flare activity, users should select a scan mode of “Solar Attenuators with High Noise Internal Cal”.

      These selections ensure that cal noise levels of these scans are of order a few percent of the anticipated system temperature. (A third scan mode, “Noise Reverse Coupler Setup (L Band only)”, is not yet supported.)

      Solar observations rely on the switched power system to flux calibrate the data. This, in turn, requires the use of 8-bit sampling modes in the frequency bands that support solar observing. Solar observing is such that input power levels may change from scan to scan, particularly for active Sun targets. This being the case each source scan requires a setup scan to trigger the system to reset the requantizer gain.

      Considerations for the calibration of solar observations are otherwise quite similar to standard gain calibration. Phase calibration proceeds in the same manner for solar observations as it does for general observing. However, it is important to be aware of two factors: First, when observing a calibrator source, the attenuation used when observing the Sun is removed from the signal path; the RFI that is suppressed, to a large degree, when observing the Sun is fully present in observations of calibrators. Second, when the attenuators are removed, solar radiation may be present in the sidelobes of the primary beam if the phase calibrator is too close to the Sun, particularly if the Sun is in an active phase. Therefore, it is advisable to observe phase calibrators that are further in angular distance from the source than would normally be used.

      VLA+LWA (eLWA) Observing

      The VLA can be used to record individual antenna signals (voltages) to VLBI Data Interchange Format (VDIF) files. Similar to Very Long Baseline Interformetry, using the phased-VLA or individual antennas, this specifically allows for joint correlation of VLA antennas with the Long Wavelength Array (LWA) stations in New Mexico in the overlapping 4m-band range. Joint correlation is performed offline using the LWA software correlator, which is located inside the VLA control building and produces FITS-IDI compatible data output.

      For shared risk observing, this mode is available for a single subband with a center frequency of 76 MHz, a bandwidth of 8 MHz, and 4-bit VDIF output. Other possible modes or center frequencies (limited by the VLA MJP-dipole response) are available through resident-shared risk observing. Use of any RSRO modes should also be brought to the attention of the LWA Director in order to verify that correlation is possible.

      eLWA proposals that are granted VLA time are automatically granted time for the LWA stations, for more information on how to select the eLWA resource for your proposal, refer to the PST manual. Observations are prepared and scheduled like any other VLA observations through the OPT and the LWA stations are automatically triggered to follow the VLA pointing. Specific instructions for eLWA instrument setups are provided in the OPT manual.

      The two LWA stations currently available in this mode are: LWA1 (close to the center of the VLA) and LWA-Sevilleta (on Sevilleta National Wildlife Refuge, ~80 km North-East of the VLA).

       

      Current limitations

      • VLA+LWA (eLWA) observing cannot be mixed with other observing modes in a single scheduling block/observing script. This limitation is due to the disabling of the array geometric delay model.
      • Data quality is highly susceptible to the low-frequency interference environment. Known sources of interference can be the active sun, powerline arcing, or self-generated interference from AC power components or digital electronics. The environment is regularly monitored and mitigation measures are taken, which in rare cases can significantly delay or prevent execution of observations.
      • Correlated data products are available through the LWA archive only and will eventually also be available through the regular NRAO archive. The PI of a proposal will be notified when correlated products are available.
      • Data calibration is recommended to be performed using AIPS.

      WIDAR Correlator

      Introduction

      The correlator configurations offered for general observing may be divided into three basic modes: wideband, spectral line, and subarrays. The possible setups are also subject to the integration time and data rate restrictions outlined in the section on Time Resolution and Data Rates. The possibilities and restrictions are embodied in the Resource Catalog Tool for proposing (RCT-proposing) and in the Resources section of the Proposal Submission Tool (PST), which must be used to define the correlator configuration for General Observing (GO) and Shared Risk Observing (SRO) proposals.

      Additionally, phased-array configurations are possible.  These are allowed for VLBI experiments (see the section on VLBI Observations) and for phased-array pulsar observations.

      Wideband and spectral line observing modes with the WIDAR correlator are described below. For the subarray mode, we refer to the Subarrays section of the OSS, and for pulsar observing modes, we refer to the Pulsar Observing section of the OSS.

      For technical details about the WIDAR correlator, refer to References 14 in Documentation.

       

      WIDAR Correlator: Wideband Observing

      The wideband observing setups provide the widest possible bandwidth for a given observing band, with channel spacing depending on the number of polarization products as listed in the following table 4.1.1:

       

      Table 4.1.1: Wideband & Subarray Correlator Options
      (all but P and L-bands)
      Polarization products Channel spacing
      Full (RR, RL, LR, LL) 2 MHz
      Dual (RR and LL) 1 MHz
      Single (RR or LL) 0.5 MHz

       

      8-bit wideband setups are available for all observing bands, providing a total of 2 GHz of bandwidth per polarization (1 GHz per polarization at L-band, and 256 MHz per polarization at P-band). 3-bit setups are available for all bands above S-band, providing total bandwidths per polarization of 4 GHz (C/X-bands), 6 GHz (Ku-band), or 8 GHz (K/Ka/Q-bands). In all cases, except for P and L-band, each of the subbands is 128 MHz wide. At L-band the default is 64 MHz/subband, yielding channels twice as narrow as those listed in the table above, while at P-band the default is 16 MHz/subband, resulting in 125 kHz channel spacing.

      In many frequency bands, the total processed bandwidth is less than that delivered by the front-end. In those cases, the observer may independently tune two 1 GHz baseband pairs when using the 8-bit samplers, or four 2 GHz baseband pairs when using the 3-bit samplers, or choose to have a mix 8-bit and 3-bit samplers. The tuning restrictions are described in the section on VLA Frequency Bands and Tunability, and the 8-bit and 3-bit samplers are described in the section on VLA Samplers.

       

      WIDAR Correlator: Spectral Line Observing



      Basebands and Subbands

      Observers have access to very flexible correlator configurations using up to 64 subbands in up to 4 basebands sampled with the 8-bit and/or the 3-bit samplers. These capabilities may be summarized as follows:

      • Two 1 GHz baseband pairs using the 8-bit samplers, or four 2 GHz baseband pairs using the 3-bit samplers, independently tunable within the limits outlined in the section on VLA Frequency Bands and Tunability. The 8-bit baseband pairs are referred to as A0/C0 and B0/D0, while the 3-bit samplers are A1/C1, A2/C2, B1/D1, and B2/D2. The AC/BD nomenclature corresponds to that of the IF pairs in the pre-expansion VLA.
      • Up to 16 subband pairs (spectral windows) in each 3-bit baseband pair, and up to 32 subbands in each 8-bit baseband pair, for a total of up to 64 subbands in any correlator configuration:
        • Tuning, bandwidth, number of polarization products, and number of channels can be selected independently for each subband;
        • All subbands must share the same integration time;
        • No part of a subband can cross a 128 MHz boundary;
        • Subband bandwidths can be 128, 64, 32, …, 0.03125 MHz (128 / 2n, n=0, 1, …, 12).
      • The sum over subbands of channels times polarization products is limited to 16384 (without recirculation):
        • These may be spread flexibly over subbands and polarization products, in multiples of 64: 64, 128, 192, 256, 384, …, 16384 cross-correlation products;
        • Recirculation may be used to increase the number of channels per subband for subbands narrower than 128 MHz. Baseline Board stacking may be used to increase the number of channels per subband for setups requiring less than 64 subbands;
        • Assigning many channels to a given subband may reduce the total bandwidth and/or the total number of subbands available.

      The remainder of this section discusses the various limitations in more detail, including some examples to show how they come up in practice.

       

      Subband Tuning Restrictions

      Each subband may be placed anywhere within a baseband, with the caveat that no subband may cross a 128 MHz boundary. Mathematically:

      νBB0 + n×128 MHz <= νsbLow <= νsbHigh <= νBB0 + (n+1)×128 MHz

      where:

      νBB0 the lower frequency edge of the baseband;
      n= 0, 1, …, 7 (, …, 15) (any integer between 0 and 7 for 8-bit, between 0 and 15 for 3-bit);
      νsbLow
      the lower edge of the subband
      (the subband center frequency minus half the subband bandwidth);
      νsbHigh
      the upper edge of the subband
      (the subband center frequency plus half the subband bandwidth).

      For example, if the baseband were tuned to cover 10000–11024 MHz, one could place a 64 MHz subband to cover 10570–10634 MHz, but not to cover 10600–10664 MHz because that would cross the 128 MHz boundary at 10640 MHz. Note in particular that the center of a baseband is a boundary and no line should be observed at the baseband center.

      The figure below illustrates these restrictions:

      Correlator configuration figure: bandpass8jul12.png

      The black curve shows the analog filter response for an 8-bit baseband covering 1024 MHz, starting at νBB0. The dashed blue vertical lines show the 128 MHz boundaries; no subband can cross those boundaries and 128 MHz subbands are thus constrained to cover a region between two of those boundaries, with no finer tuning being possible. Narrower subbands, like the 64 MHz subband shown here in red, can be shifted around arbitrarily within one of the 128 MHz slots, but cannot cross any of these boundaries. (The dotted vertical red lines show the boundaries of the 64 MHz subband, while the solid curve shows an illustrative line within the subband.)

      The analog filter shape defining the baseband rolls off severely at one edge of the baseband, so the 128 MHz slot at that edge has reduced sensitivity. The baseband edge is at the lowest sky frequency in the baseband when using the upper sideband, and at the highest sky frequency in the baseband when using the lower sideband.

       

      Subband Bandwidths and the Digital Filter Response

      The bandwidth for each subband may be selected independently, and can be any of 128/2n MHz, for n= 0, 1, …, 12: 128, 64, 32, 16, 8, 4, 2, or 1 MHz, or 500, 250, 125, 62.5, or 31.25 kHz.

      The usable portion of the subband is set by three effects. First, as discussed above, the analog filters which define the baseband are not perfect, leading to lower sensitivity in the 128 MHz near the baseband edge for the 8-bit samplers.

      Second, because the digital filters are not infinitely sharp, the rejected sideband leaks in at both edges of the subband. This leads to additional (aliased) noise, with a factor ~2 increase in the noise at the subband edges, dropping to a few percent within a few percent of the subband edge. The precise filter shape and noise increase is a complex but predictable function of the subband bandwidth (sbBW) and the subband tuning.

      The third effect stems from the offset frequencies used for sideband rejection in the WIDAR correlator. The local oscillators at the individual antennas are tuned to slightly different frequencies, with those offsets taken out in the correlator. This means that each antenna observes a slightly different sky frequency, and thus some baselines will not give an interesting correlation near one edge of the subband. The maximum frequency shift is currently set to 32×f0, with the fundamental f0 being set to f0 = max(25.6 kHz×sbBW/128 MHz, 100 Hz). Here sbBW is the smallest subband bandwidth within the baseband. For the wider subband bandwidths the maximum frequency shift corresponds to <1% of that bandwidth, but for narrower subbands the effect can be severe. For instance, a 31.25 kHz subband has f0 = 100 Hz, and a maximum frequency shift of 3.2 kHz—10% of the subband may be lost on some baselines.

       

      Spectral Channels and Polarization Products

      Each subband, without recirculation enabled, can have a different number of channels and polarization products, subject to two limitations:

      1. For the ithsubband, the number of spectral channels can be:
        • 64 nBlBP,i with full polarization products (RR,RL,LR,LL)
        • 128 nBlBP,i with dual polarization products (RR and LL)
        • 256 nBlBP,i with a single polarization product (RR or LL)
        Here nBlBP,i= 1, 2, 3, 4, 5, …, 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.
      2. The sum over all subbands of nBlBP,i must be less than or equal to 64, the number of Baseline Board pairs in the correlator. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64 × 256 = 16,384 (without recirculation).

      Baseline Boards are the boards in the WIDAR correlator where the actual cross-multiplications are done. There are 128 Baseline Boards arranged as 64 Baseline Board pairs (BlBPs). The limitations given here correspond to the capabilities of the individual boards and the finite number of boards the correlator has. Use of more than one pair per subband (i.e., nBlBP>1) is known as Baseline Board stacking; see additional details about this below.

      Limitation #1 corresponds to table 4.2.1 of the options for subband bandwidth and spectral resolution when using nBlBP Baseline Board pairs for a subband:

      Table 4.2.1: Subband Bandwidth and Spectral Resolution Options (without recirculation)
      Subband bandwidth &
      total velocity coverage
      Full polarization products
      (RR, RL, LR, LL)
      64nBlBP spectral channels

      Channel spacing:
      Dual polarization products
      (RR and LL)
      128nBlBP spectral channels
      Channel spacing:
      Single polarization product
      (RR or LL)
      256nBlBP spectral channels

      Channel spacing:
      128 MHz 38400/νGHz km/s 2000/nBlBP kHz 600/nBlBPGHz km/s 1000/nBlBP kHz 300/nBlBPGHz km/s 500/nBlBP kHz 150/nBlBPGHz km/s
      64 19200 1000 / nBlBP 300 / nBlBP 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP
      32 9600 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP
      16 4800 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP
      8 2400 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP
      4 1200 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 /n BlBP
      2 600 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP
      1 300 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP
      0.5 150 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP
      0.25 75 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP
      0.125 37.5 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP
      0.0625 18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
      0.0325 9.375 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP 0.122 / nBlBP 0.037 / nBlBP

      Subband bandwidth and spectral resolution options. Note that the table entries refer to the spacing between spectral channels—that spacing is before any frequency smoothing, so these channels are not independent.

      • nBlBP is the number of Baseline Board Pairs assigned to the subband.
      • Each subband may have a different number of spectral channels and polarization products, and each may be tuned independently.
      • There can be at most 16 subbands per baseband, and nBlBP must be an integer: 1, 2, 3, 4, 5, …, 64.
      • The sum of nBlBP over all subbands must be less than or equal to 64.
      • Use of more than one BlBP for a subband may further restrict the number of subbands available in one or more of the basebands; see text for details.

      Here are four examples of allowed general observing setups which use all 64 BlBPs to produce the maximum number of channels times polarization products:

      Table 4.2.2: Example BlBP Setups
      Baseband Subband

      Pol'n
      Products

      Spectral
      channels

      nBlBP
      Example 1 A0/C0 sb0 RR 16384 64
      Example 2 A0/C0 sb0 RR 8192 32
      A0/C0 sb1 RR, LL 1024 8
      A0/C0 sb2 RR, LL 512 4
      B0/D0 sb0 RR, LL 2048 16
      B0/D0 sb1 RR,RL,LR,LL 256 4
      Example 3 A0/C0 sb0 RR 8192 32
      A0/C0 sb1 LL 1024 4
      A0/C0 sb2 RR, LL 1024 8
      A0/C0 sb3 RR,RL,LR,LL 1024 16
      A0/C0 sb4 RR,RL,LR,LL 256 4
      Example 4 A0/C0 sb0-5 RR,RL,LR,LL 64 6 x 1
      A0/C0 sb6 RR, LL 3840 1 x 30
      A0/C0 sb7 RR 768 1 x 3
      A0/C0 sb8 RR,RL,LR,LL 192 1 x 3
      B0/D0 sb0-2 RR,RL,LR,LL 64 3 x 1
      B0/D0 sb3 LL 768 1 x 3
      B0/D0 sb4 RR, LL 2048 1 x 16

       

      Recirculation

      When the subband bandwidth is less than the maximum 128 MHz, the spare clock cycles that become available in the correlator hardware can be re-purposed to compute additional lags using a single baseline board. This increases the spectral resolution within the subband, and is known as recirculation. Each factor of two reduction in subband bandwidth results in an additional factor of two maximum lags; therefore for subbands of 128 MHz / N the possible spectral resolution (in units of frequency) can be increased by a factor of N2.

      Recirculation vs. Baseline Board Stacking

      When faced with the choice between recirculation and Baseline Board stacking to increase the number of channels in a subband, we recommend recirculation for subbands narrower than 128 MHz; this is supported in observatory software (RCT-proposing, OPT). Using recirculation rather than stacking frees up more Baseline Boards for other uses; alternatively the experiment becomes less dependent on all Baseline Board pairs being available/working at the time of observation. For subbands of 128 MHz, recirculation is not possible, and Baseline Board stacking must be utilized to increase the number of channels.

      The present implementation of recirculation is that, for each halving of the subband bandwidth, the number of channels in the subband may be doubled without having to use additional correlator hardware. The maximum recirculation factor for a subband is 128/(subband bandwidth in MHz) and, of course, subject to other configuration restrictions such as data rate.

      Baseline Board Stacking

      As opposed to recirculation, which increases the number of channels in a subband by exploiting otherwise unused CPU resources, Baseline Board stacking adds more channels to a subband by adding correlator hardware resources, i.e., using up more Baseline Board pairs. Using Baseline Board stacking may therefore limit the number of subbands available in one or more of the basebands. Understanding how this works requires understanding some of the details of the correlator hardware. That understanding is built into the RCT-proposing, and observers may simply use that tool to find out whether their particular setup will, in fact, work. But the results can be confusing without some understanding of the hardware constraints from which they arise. These hardware constraints are complex, and most observers will not need to understand these details. The following section is for those who are attempting complex line experiments and who find the RCT-proposing restricting the number of subbands and/or channels they can use in unexpected ways. Most observers can skip the following section.

      Baseline Board Stacking and Correlator Use

      First let us consider how the correlator hardware is organized. The cross-multiplications in the WIDAR correlator are spread across 64 Baseline Board pairs (BlBP), arranged into 4 quadrants of 16 BlBP each. Each baseband is connected directly to one of those quadrants. In the simplest mode, each of the 16 BlBP of a quadrant handles the correlations for one of the 16 subbands of the corresponding baseband. Four basebands and four quadrants are required to handle the full 8 GHz of bandwidth per polarization provided by the 3-bit (wideband) samplers: 8 GHz is split into four basebands of 2 GHz each, with each baseband fed into a different BlBP quadrant. Each BlBP in that quadrant handles a subband of maximum bandwidth 128 MHz, so 16 BlBP handles 16 subbands for a total of 16×128 MHz = 2048 MHz.

      A single BlBP produces 256 cross-correlations per baseline for a single subband, which can be used for a single polarization product (e.g., RR or LL with 256 spectral channels), or two (RR and LL with 128 spectral channels each), or four (RR, RL, LR, and LL with 64 spectral channels each).

       

      When using the 8-bit samplers, the total bandwidth is only 2 GHz per polarization, split into two basebands of 1 GHz each. The simplest continuum setup uses only two quadrants, since there are only two basebands; and only 8 subbands are required to span the 8×128 MHz = 1024 MHz of each baseband. Three-quarters of the correlator BlBP hardware remain unused.

       

      The spectral line mode allows access to these extra correlator resources through Baseline Board stacking: using multiple BlBPs to process the same subband and produce more cross-correlations for that subband. This is done using crossbar switches which make the data for a single subband available to several BlBPs. Those BlBPs can then be used to produce more spectral channels for that subband, with n BlBPs producing 256×n cross-correlations per baseline. The limit on the total number of cross-correlations (16384) stems from the total number of BlBPs (64): 64×256 = 16384.

      Unfortunately, completely flexible crossbar switches are expensive and could not be implemented in the VLA's correlator. This means that one cannot route a given subband to a randomly-chosen BlBP. The routings which are possible, are as follows:

      1. A subband in a baseband can be routed to any BlBP within the corresponding quadrant.
      2. Data coming into a given BlBP in one quadrant, can be routed to the corresponding BlBP in any other quadrant.

      Routing option #1 means that one could use all the BlBPs within a quadrant to correlate a single subband, yielding 16×256 = 4096 cross-correlations for that subband:

       

      Routing option #2 means that one could use the BlBPs in all 4 quadrants to correlate a single subband. One simple case would use 4 BlBPs to correlate each of the 16 subbands in a single baseband, yielding 4×256 = 1024 cross-correlations for each of those subbands. Note that in this case, no BlBPs are left to correlate any data from the second baseband.

      Using routing option #2 does come with a subtle cost: assigning a BlBP in quadrant X to correlate a subband corresponding to quadrant Y removes that BlBP from use in the baseband corresponding to quadrant X…and therefore also removes the corresponding subband in that baseband. So, getting more channels for a subband in one baseband may prevent the use of a subband in a different baseband. To take a simple example, consider an experiment where one wishes to observe a single line in dual polarization with 512 channels (requiring 4 BlBPs), plus as much continuum bandwidth as possible. Naively, one would say there are 16 subbands in each baseband; one is used for the spectral line, so that leaves 16+15 = 31 subbands, and with the widest subband bandwidth (128 MHz) the total available continuum should be 31×128 MHz = 3968 MHz per polarization. Actually, however, there are only 15+15 subbands available, or 30×128 MHz = 3840 MHz per polarization, because the spectral line subband has eaten one BlBP corresponding to the other baseband:

      If the same spectral line required twice as many channels, this will result in the loss of two subbands in both of the basebands:

      In some cases one may want to use a different routing to use up subbands in one baseband in preference to another. For instance, the same spectral line setup (2048 cross-correlations for a single spectral line subband, plus as much continuum as possible) could be set up to allow 13 continuum subbands in the A0/C0 baseband, and the full 16 continuum subbands in B0/D0:

      Understanding these confusing constraints can help observers set up the VLA more effectively to achieve their scientific goals. For instance, in a mixed line+continuum experiment, it works best to use the resource tools to set up the baseband tunings and subband channelization for the most important lines first, then add the desired continuum, and then see what correlator resources remain for any lines of secondary interest.

      The above examples all use BlBP pair stacking in powers of 2, but this is not required. To give some idea of more complex possibilities, the following tables (4.2.3 and 4.2.4) give two examples of other possible configurations. The RCT display shows how the Baseline Boards are used to process the individual subbands. The cyan boxes (shaded when printed out in black and white) show the Baseline Boards used to process data from baseband A0/C0, while the yellow boxes show Baseline Boards used to process data from baseband B0/D0.

      Table 4.2.3: Complex Configuration Example #1
      Baseband Subband Pol'n products Spectral channels nBlBP
      A0/C0 sb0 RR 10240 40
      A0/C0 sb1 LL 768 3
      A0/C0 sb2 RR,LL 2176 17
      B0/D0 sb0 RR 256 1
      B0/D0 sb1 RR,LL 384 3
      RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

       

      Table 4.2.4: Complex Configuration Example #2
      Baseband Subband Pol'n products Spectral channels nBlBP
      A0/C0 sb0 RR 4352 17
      A0/C0 sb1 RR, LL 1152 9
      B0/D0 sb0 RR,RL,LR,LL 192 3
      B0/D0 sb1 RR, LL 4480 35
      RCT display: corr-cfg-fig:sro2_8bit_ac17+9_bd3+35

       

      The individual subbands can have different bandwidths, and those bandwidths may be chosen completely independently of the number of spectral channels in each subband. So, for instance, a subband with a bandwidth of 2 MHz and 1152 spectral channels would have a channel separation of 2 MHz/1152 = 1.736 kHz; but the observer could equally well choose a bandwidth of 64 MHz for that subband, leading to a channel separation of 64 MHz/1152 = 55.56 kHz.

      Use of the 3-bit samplers further extends the possibilities. Here is one example:

      Table 4.2.5: 3-bit Complex Configuration Example #1
      Baseband Subband Pol'n products Spectral channels nBlBP Quadrant(s): Column(s)
      A1/C1 sb0-8 RR, LL, RL, LR 9 x 64 9 x 1 Q1: 0–8
      A1/C1 sb9 RR, LL 1 x 1152 1 x 9 Q1 & Q3: 9–11, 14 / Q4: 9
      A1/C1 sb10 RR 1 x 1792 1 x 7 Q1 & Q3 & Q4 : 12,13 / Q2: 13
      A1/C1 sb11 RR, LL 1 x 384 1 x 3 Q1 & Q2 & Q3: 15
      A2/C2 sb0-11 RR, LL, RL, LR 12 x 64 12 x 1 Q2: 0–11
      A2/C2 sb12 LL 1 x 768 1 x 3 Q2: 12, 14 / Q4: 14
      B1/D1 sb0-3 RR, LL, RL, LR 4 x 64 4 x 1 Q3: 0–3
      B1/D1 sb4 RR, LL, RL, LR 1 x 320 1 x 5 Q3: 4–8
      B2/D2 sb0-6 RR, LL, RL, LR 7 x 64 7 x 1 Q4: 0–6
      B2/D2 sb7 RR, LL 1 x 640 1 x 5 Q4: 7, 8, 10, 11, 15
      RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

       

      Once again, the RCT-proposing implements all of these constraints.

      Documentation

      Documentation

      Documentation for VLA data reduction, image making, observing preparation, etc., can be found in various manuals. Current manuals are available on-line. Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the VLA user, not listed here, are available from our website.

      1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: "Synthesis Imaging in Radio Astronomy." This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. The lectures of the 2014 workshop are available at the 14th Synthesis Imaging Workshop web site.
      2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series.
      3. GUIDE TO OBSERVING WITH THE VLA: Describes details of how to observe with the VLA once you have been allocated time on the VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide). Including special observing modes such as:
        1. CALIBRATION (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/calibration)
        2. OBSERVING WITH THE 8-BIT (up to 2 GHz bandwidth) & 3-BIT (up to 8 GHz bandwidth) SAMPLER SYSTEMS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up);
        3. SPECTRAL LINE OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/line); 
        4. HIGH FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/hifreq);
        5. LOW FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/lofreq);
        6. VERY LOW FREQUENCY OBSERVING (< 500 MHz) (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/vlofreq);
        7. POLARIMETRY (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/pol);
        8. MOSAICKING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/mosaicking);
        9. RADIO FREQUENCY INTERFERENCE (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/rfi);
        10. MOVING OBJECTS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/moving);
        11. VLBI AT THE VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/vlbi).
      4. *CASA COOKBOOK (deprecated with last updates for CASA 4.7.2): The CASA Cookbook for use of the package for data reduction of VLA (& ALMA) data is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See (http://casa.nrao.edu/docs/cookbook/)
      5. CASA Online Documentation: https://casadocs.readthedocs.io/en/stable/
      6. VLA CASA Guides: Tutorials and data reduction examples of VLA data in CASA (https://casaguides.nrao.edu/index.php/Karl_G._Jansky_VLA_Tutorials)
      7. *AIPS COOKBOOK: The Astronomical Image Processing System (AIPS) software is able to fully calibrate VLA data and do most imaging operations. The exception is the wide-band (bandwidth synthesis) deconvolution which is being developed in CASA only. ALMA data may also be reduced in AIPS although the package is not fully qualified to calibrate data from the ALMA linearly-polarized feeds. The Cookbook description for calibration and imaging under the AIPS system can be found near all public workstations in the SOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See (http://www.aips.nrao.edu/cook.html)
      8. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See (http://www.aips.nrao.edu/goaips.html)
      9. *VLA CALIBRATOR LIST: This page contains the list of VLA Calibrators in both 1950 and J2000 epoch. See (https://science.nrao.edu/facilities/vla/observing/callist)
      10. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983.
      11. *HISTORICAL VLA MEMO SERIES: archive memo series from the early days of the VLA. See (http://library.nrao.edu/vlam.shtml)
      12. *RECENT VLA MEMO SERIES: the memo series relating to the expanded capabilities of the VLA. See (http://library.nrao.edu/evla.shtml)
      13. *The VLA Expansion Project: Construction Project Book. The Expanded VLA Project Books contains the technical details of the VLA Expansion construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml.
      14. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains a useful source of information on much of the VLA. There are a few hard copies at the VLA and in the DSOC. Much of this document is now available for download (https://science.nrao.edu/facilities/vla/obsolete/green-book). Note: it does not include any information about the hardware and software specific to the expanded Karl G. Jansky VLA.
      15. WIDAR: The DRAO design and development documents of the WIDAR correlator of the VLA are available at http://www.aoc.nrao.edu/widar/docs/.

       

      Online Tools & Important Links

      The NRAO User Portal. (https://my.nrao.edu) This is a gateway to the NRAO interactive services that include the Proposal Submission Tool (PST).

      The NRAO Proposal Submission Tool (PST) online manual. (https://science.nrao.edu/facilities/vla/docs/manuals/proposal-guide/pst)

      The VLA Exposure Calculator Tool (ECT) online manual. (https://science.nrao.edu/facilities/vla/docs/manuals/propvla/determining)

      The VLA Exposure Calculator Tool (ECT). (https://obs.vla.nrao.edu/ect/)

      The Resource Catalog Tool for proposers. (https://rctp.vla.nrao.edu/rct/)

       

      Acknowledgements

      Many thanks to all the VLA staff and our RSRO participants who have worked long and hard to commission these capabilities and who have helped to create this extensively updated set of documentation.

      NRAO is grateful to Professor Rob Ivison for supporting the upgrade of some of the 3-bit samplers on the VLA via a grant from the European Research Council. For observations using the 3-bit samplers between May 2015 and March 2018 we encourage users to include the following text in the Acknowledgments section of their publications:

      "We acknowledge funding towards the 3-bit samplers used in this work from ERC Advanced Grant 321302, COSMICISM."

      Contact Information

      Please go to the People page for information on key personnel at NRAO-Socorro.

      Please direct queries to the NRAO Helpdesk; you can expect a response within one to two business days.  


       

      Editor's Notes

      This Observational Status Summary for the Karl G. Jansky (expanded) VLA is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. For questions on the content, or suggestions that would enhance the clarity of this guide, we recommend contacting the NRAO Helpdesk.

      files

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      test collection

      VLA Observational Status Summary 2016A

      VLA capabilities February 2016 - October 2016

      Introduction

      Purpose of Document, Older Versions of the OSS

      This Observational Status Summary (OSS) summarizes the instrumental status of the Karl G. Jansky Very Large Array (VLA) for the C, CnB, and B configurations for the observing period 5 February 2016 through 3 October 2016 (2016A semester), and should be used when preparing proposals for the 3 August 2015 deadline.  For capabilities prior to that date, we refer to our overview of all OSS versions available online.

      The OSS is intended as a ready reference for those contemplating use of the VLA for their astronomical research. The information is in summary form - those requiring greater detail should use our Helpdesk, or refer to the manuals and documentation listed in Documentation. Most of the information contained here, and much more, is available through the VLA science web pages, and the companion document for the VLBA.

      The VLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for efficient use to be made of it.  Although the NRAO strives to make using the VLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first-time users are encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. Refer to the Visiting the DSOC and VLA page for details. The VLA is an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it.

      An Overview of the VLA

      The Karl G. Jansky Very Large Array (VLA) is a 27-element interferometric array, arranged along the arms of an upside-down "Y", which will produce images of the radio sky at a wide range of frequencies and resolutions. It is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (SOC) in Socorro, New Mexico.

      The basic data produced by the VLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array's elements.  The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency.   Another mode of observing (commonly called phased array) allows operation of the array as a single element through coherent summation of the individual antenna signals. This mode is most commonly used for Very Long Baseline Interferometry (VLBI) observing and for observations of rapidly varying objects, such as pulsars.

      The VLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic antenna arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A, respectively. In addition, there are 3 "hybrid" configurations labelled DnC, CnB, and BnA, in which the North arm antennas are deployed in the next larger configuration than the SE and SW arm antennas. These hybrid configurations are especially well suited for observations of sources south of δ = −15° or north of δ = +75°, for which the foreshortening of the longer North arm results in a more circular point spread function.  For details about antenna positions in the various configurations we refer to the stations position file (ps version, pdf version).

      The VLA completes one cycle through all four configurations in an approximately 16 month period. Consult the Configuration Plans and Proposal Deadlines page or recent NRAO and AAS newsletters for current and up-to-date configuration schedules and associated proposal deadlines. Refer to the Guide to Proposing for the VLA for information on how to submit an observing proposal.

      Observing projects on the VLA will vary in duration from as short as 1/2 hour to as long as several weeks.  Most observing runs have durations of a few to 24 hours, with only one, or perhaps a few, target sources.  However, since the VLA is a two-dimensional array, images can be made with data durations of less than one minute.  This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects.  See Snapshots for details.

      All VLA antennas are outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers cover the frequency ranges of 1-2 GHz, 2-4 GHz, 4-8 GHz, 8-12 GHz, 12-18 GHz, 18-26.5 GHz, 26.5-40 GHz, and 40-50 GHz and the bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively.  In addition, all antennas of the VLA have receivers for lower frequencies, enabling observations at P-band (230-470 MHz). These low frequency receivers also work at 4-band (54 to 86 MHz), and new feeds have been deployed on a small number of VLA antennas to observe at this frequency range.

      The VLA correlator is both powerful and flexible.  Details of the correlator configurations being offered for VLA science are described in the WIDAR Section of this document.   It is important to realize that the VLA correlator is fundamentally a spectral line correlator and that even 'continuum' observations are done in a wide-band mode with many channels.

      The Expanded Very Large Array (EVLA) Project

      The Expanded VLA (EVLA) project modernized the VLA electronics (built in the 1970s and 1980s) in order to improve several key observational parameters by an order of magnitude or more.  Some of the details of the EVLA Project may be found at http://www.aoc.nrao.edu/evla/.   The EVLA project was funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico.   Total funding was approximately USD 94 million in Year 2006 dollars, including USD 59 million in new NSF funding, USD 16 million in redistributed effort from the NRAO Operations budget, USD 17 million for the correlator from Canada, and USD 2 million from Mexico.

      The EVLA project was completed on time and on budget at the end of 2012, 11 years after it began.  Its key observational goals were (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA's 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator capable of processing the large bandwidth with a minimum of 16,384 spectral channels per baseline.  All goals were met.  A comparison of some of the EVLA performance parameters with those of the original VLA is provided in Table 1.

      Note: The "Factor" gives the factor by which the EVLA parameter improves on the equivalent VLA parameter.

      VLA to EVLA transition

      The correlator that was the heart of the VLA for three decades was decommissioned on 11 January, 2010, and replaced with the new EVLA "WIDAR" correlator.  The VLA was shut down to outside users until March 2010, during which time hardware was transferred from the old correlator to the EVLA correlator and observing modes commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles also changed, from ABCDA to DCBAD, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. The last VLA antenna was retrofitted to EVLA specifications in May 2010.

      The first full configuration cycle of early science using the EVLA correlator saw up to 256 MHz of bandwidth offered to the general community, and 2 GHz bandwidth for observers willing to visit Socorro to help with EVLA commissioning.  By the end of 2011, at the start of the second configuration cycle, up to 2 GHz bandwidth was offered to the general community.  This increased to 8 GHz at the start of full operations in semester 2013A.

      Proposing for the VLA for the Next Semester

      The 2016A Call for Proposals

      The 2016A Call for Proposals details the General Observing (GO) capabilities being offered for the Karl G. Jansky Very Large Array (VLA).

      In addition to these general capabilities, NRAO continues to offer shared risk observing options for those who would like to push the capabilities of the VLA beyond those offered for general use.  These are the "Shared Risk Observing" (SRO) and "Resident Shared Risk Observing" (RSRO) programs.

      Details about what is being offered for each program along with a list of the steps that should be completed for each proposal category is given below.  In addition, a summary of the links to updated documentation and tools to prepare your proposal is also given below, including the Guide to Observing with the VLA.

      If you have any questions or problems with any link or tool, please submit a ticket through the NRAO Helpdesk.

      Note: The observing semester 2016A is the last semester NRAO will offer the hybrid configuration, and only the CnB hybrid may be requested.

      Considering the lack of future hybrid configurations (after semester 2016A), guidelines on how to substitute such configurations with the use of principal array configurations are presented in the Array Configurations section of the Guide to Proposing for the VLA.

       

      General Observing (GO) and Shared-Risk Observing (SRO)

      Summary of Capabilities

      As described in the Call for Proposals, the VLA offers continuous frequency coverage from 1-50 GHz in the following observing bands: 1-2 GHz (L-band); 2-4 GHz (S-band); 4-8 GHz (C); 8-12 GHz (X); 12-18 GHz (Ku); 18-26.5 GHz (K); 26.5-40 GHz (Ka); and 40-50 GHz (Q).  Both single pointing and mosaics with discrete, multiple field centers will be supported under General Observing (GO).  P-band (230-470 MHz) Stokes I continuum observations are also offered under GO. See the Low Frequency Observing section below for more details. Data rates of up to 60 MB/s (216 GB/hour) will be available to all users, combined with correlator integration time limits per band and per configuration, as described in the Time Resolution and Data Rates. However, data rates in excess of 25 MB/s (90 GB/hour) require additional justification.  Limitations on frequency settings and tuning ranges are described in the Frequency Bands and Tunability Section.

      The GO capabilities being offered are:

       

      Capability Description
      8-bit samplers
      • Standard default set-ups for:
        • 2 GHz bandwidth continuum observations at S/C/X/Ku/K/Ka/Q bands (16 x 128 MHz sub-bands)
        • 1 GHz bandwidth continuum observations  at L band (16 x 64 MHz sub-bands)
        • 256 MHz bandwidth continuum observations at P band (16 x 16 MHz sub-bands)
      • Flexible set-ups for spectroscopy, using two independently tunable 1 GHz baseband pairs, each of which can be split into up to 16 flexibly tunable sub-bands
      • Single, dual, and full polarization products
      3-bit samplers
      • Standard default set-ups for:
        • 8 GHz bandwidth continuum observations at K/Ka/Q bands
        • 6 GHz bandwidth at Ku band
        • 4 GHz bandwidth at C/X bands
      • Flexible set-ups for spectroscopy, using four independently tunable 2 GHz baseband pairs, each of which can be split into up to 16 flexibly tunable sub-bands
      • Single, dual, and full polarization products
      Mixed 3-bit and 8-bit samplers
      • Allows more flexibility for simultaneous continuum and high-resolution spectral line observing

      Sub-arrays

      • Up to 3 independent sub-arrays using standard 8-bit continuum set-ups
      Phased-array for VLBI
      • VLA Phased Array ("Y27") for VLBI (see the VLBA call for proposals and the VLBA Observational Status Summary for a detailed description of the VLA Phased Array capabilities being offered)

       

      SRO capabilities can be set up via the Observing Preparation Tool (OPT) and run through the dynamic scheduler (without intervention) but are not as well tested as GO capabilities.  A summary of the SRO capabilities being offered are:

          • On-the-Fly (OTF) mosaicing
          • Up to 32 sub-bands per baseband with the 8-bit samplers
          • recirculation of up to a factor of 64
          • 8-stream recording with phased VLA for VLBA observing.

        ‡Note: OTF mosaic observations have certain limitations on the length and number of scans. Contact the VLA staff through the NRAO helpdesk to check if the OTF mosaic observations to be proposed are technically feasible ahead of submitting a proposal for the 2016A semester.

        We expect that most SRO programs will have no or only minor problems that can be corrected quickly.  However, if an SRO program fails and it becomes clear that detailed testing with additional expertise is needed, then the project must make an experienced member from their team available to help troubleshoot the problem.  In some cases, this may require presence of that experienced member in Socorro.  If adequate support from the project is not given, then the time on the telescope will be forfeited. The additional effort is to be determined based on discussions with the NRAO staff and management and the project team.

        Steps to Completion

        1. Read the 2016A NRAO Call for Proposals description for summaries of the capabilities being offered for GO and SRO.
        2. If you are proposing for spectral line observations, run the GO Setup Tool (GOST, a java application) to define your correlator set-up and save a snapshot of the tool GUI to disk.  Note:
      • Write a scientific justification for your proposal.  Note that technical information should be included in the Technical Justification section form and does not need to be included with the scientific justification.  Save your scientific justification as a PDF file.
      • Log into the NRAO Interactive Services page (my.nrao.edu) and click on the "Proposals" tab in the top left to create a new proposal for the VLA.  Once you are in the tool, extensive help is available in the tool by clicking on the "Help" button in the top right of the tool interface (see also the links to the help documentation below).  In the PST:
        • Fill in the relevant fields for each section of your proposal;
          • If you are proposing for continuum observations, select Observing Type = "Continuum" in the General section of your proposal.  Then add a Continuum Resource, selecting the appropriate band you want to observe in.  Default optimum continuum setups for each band are defined in the PST.
          • If you are proposing for spectral line observations, Select Observing Type = "Spectroscopy" in the General section of your proposal.  Then add a Spectroscopic Resource and attach the GOST snapshot (that you saved previously) to this resource.  The GOST snapshot describes everything needed to specify the lines you want to observe and what correlator resources are needed for the observation.
        • Upload your scientific justification as a PDF file;
        • Technical justifications are now included as a separate section in the proposal.  Click on the Technical Justification page in the proposal and fill in the appropriate boxes (example of technical justification). In order to determine sensitivities, you will need to use the VLA Exposure Calculator Tool (ECT) to understand what sensitivity you will get for a given frequency, bandwidth, and integration time (guidelines for running the ECT). If you have trouble running the ECT please contact the NRAO Helpdesk;
        • Note that the PST allows observers to specify multiple resource types (e.g., you can have one proposal that specifies general and/or RSRO correlator resources).  If any resource is RSRO, it will be a RSRO proposal.
      • When your proposal is complete, validate it to make sure there are no obvious omissions or mistakes.
      • When you are satisfied, submit your proposal in the PST. Note that proposals can be withdrawn and resubmitted before the deadline.
      •  

        Resident Shared Risk Observing (RSRO)

        Summary of Capabilities

        The RSRO program provides access to extended capabilities of the VLA that require additional testing in exchange for a period of residence to help commission those capabilities. Capabilities that would fall under the RSRO program include:

          • correlator dump times shorter than 50 msec, including integration times as short as 5 msec for transient detection;
          • pulsar observations;
          • data rates above 60 MB/s;
          • recirculation beyond a factor of 64 in the correlator;
          • P-band system (230 to 470 MHz) polarimetry and spectroscopy;
          • 4-band system (54 to 86 MHz; see the Low Frequency Observing section below);
          • more than 3 sub-arrays, or sub-arrays with the 3-bit system;
          • complex phased array observations (e.g., pulsar and complex VLBI observing modes); and
          • frequency averaging in the correlator: a new capability for averaging to wider frequency channels in the correlator has been developed that will reduce the data volume for all continuum subbands

         

          Steps to Completion

          1. Read the 2016A NRAO Call for Proposals for a summary of the capabilities being offered for GO and SRO. If you want more than what is offered for GO or SRO then you are requesting a RSRO capability (one that is not well-tested or may even need additional development).  If you propose for a RSRO capability, you (or an experienced person on your team) must be able to participate in the RSRO program by coming to Socorro to help with the development and testing process.
          2. Write a scientific justification for your proposal.  Note that technical information should be included in the Technical Justification section and does not need to be included with the scientific justification.  Save your scientific justification as a PDF file.
          3. Log into the NRAO Interactive Services page (my.nrao.edu) and click on the "Proposals" tab in the top left to create a new proposal for the VLA.  Once you are in the tool, extensive help is available in the tool by clicking on the "Help" button in the top right of the tool interface (see also the links to the help documentation below).  In the PST:
          • Fill in the relevant fields for each section of your proposal;
          • Even as a RSRO proposal, you will need to create a "Resource" in the PST and select the "WIDAR RSRO" Back End.  This will give you a text field in which you can type a description of your set-up;
          • Upload your scientific justification as a PDF file;
          • Technical justifications are now included as a separate section in the proposal.  Click on the Technical Justification page in the proposal and fill in the appropriate boxes (example of technical justification). Describe what you want to do with the correlator and why this is in the RSRO Category in this section. Contact NRAO staff if you have questions about exactly what is feasible. The last box at the bottom of the Technical Justification page should be used to describe your RSRO effort. Identify who in your team can come to Socorro to help commission this capability and how their background and expertise can be applied to this development.  Working with NRAO staff, estimate the level of effort that is likely to be needed for this development and specify how long a member of your team can come to NRAO.
          • Note that the PST allows observers to specify multiple resource types (e.g., you can have one proposal that specifies general, shared-risk and RSRO correlator resources).  If any resource is RSRO, it will be a RSRO proposal.
        • When your proposal is complete, validate it to make sure there are no obvious omissions or mistakes.
        • When you are satisfied, submit your proposal in the PST. Note that proposals can be withdrawn and resubmitted before the deadline.
        • Considerations

          A RSRO proposal should contain:

          1. A scientific justification, to be peer reviewed as part of NRAO's current time allocation process, submitted through the Proposal Submission Tool.  Note that RSRO correlator resources should be specified as plain text on the Resources page in the PST by selecting "WIDAR RSRO" as the backend.
          2. The technical justification should identify the personnel who will be involved in the residency and describe how their expertise will be used to address the critical priorities of VLA development relating to their proposal.  The proposed dates of the residency must be included, so that the residency can be matched to VLA development planning.  This section will be reviewed by NRAO staff.  Limited support for accommodation in the NRAO Guest House for participants in the RSRO program may be available.

          The acceptance of a RSRO proposal will depend on the outcome of the time allocation process, and proposals will also be evaluated by NRAO staff in terms of the priorities and benefits to the VLA development and commissioning activities.

          In general, one month of resident commissioning effort is expected for every 20 hours of VLA time awarded to a RSRO project, subject to negotiation.  There is no minimum requirement for the amount of residency at NRAO.  However, the amount of time spent at NRAO to help develop the program should be realistically matched to the expected effort, including time to become familiar with relevant technical aspects.  The time proposed at NRAO should be discussed with NRAO staff to determine what is reasonable.  The length of time a RSRO expert should be needed at NRAO may be on the order of a few months.

          The period(s) of residency may occur in advance of the observing time awarded in order to decouple essential scientific requirements (such as array configuration) from other factors which may affect when personnel are available (such as teaching schedules).  However, observers should be present for one week prior to their observations in order to become familiar with the latest developments and to set up their observations.  In the special case of Target of Opportunity proposals a VLA staff collaborator may be required for setting up observations on short timescales.

          It is possible for a member of the NRAO scientific staff to satisfy the residency requirement on a RSRO proposal.  NRAO staff considering providing the residency requirement for an RSRO proposal should consult with their supervisor for further information.  Graduate students may satisfy the residency requirement, provided relevant expertise is demonstrated in the RSRO relevant sections of the proposal.  Graduate students should be accompanied by their advisor at the start of their residency.  Resident personnel will work under NRAO management in order to optimize the overall effort.  A set of clear goals will be agreed upon in advance of the start of the residency.

          The types of proposals considered under the RSRO program may include both large (>200 hours) and small (~10-200 hours) projects.  Qualified large projects proposed by consortia will be considered as long as the residency requirements are met.  A single individual may satisfy the residency requirement for several small projects.

          RSRO participation without a science proposal:

          In some cases, an individual may want to participate in development activities without writing a science proposal.  A participant may arrange to visit Socorro to contribute to development activities by submitting a proposal directly to the Assistant Director for NM Operations (nraonmad@nrao.edu) containing a proposal for who will come and the technical development to be undertaken.  If the Assistant Director approves of the request then the individual may come to Socorro to contribute to development activities.  The participant may then obtain observing time either by submitting a proposal at a regular proposal deadline, or by submitting an Exploratory Proposal through Director's Discretionary Time.  Such visits should conform to the residency requirements above.  Proposals to visit Socorro under this program may be submitted at any time.

          Low Frequency Observing (P-band and 4-band), and VLITE

          The new low frequency receiver system developed in collaboration with the Naval Research Laboratory will be available for Stokes I continuum observations at P-band (230 to 470 MHz) through the GO program. Use of the P-band system for polarimetry and/or spectroscopy will be through the RSRO Program as noted above. The new receivers also work at 4-band (54 to 86 MHz), and new feeds have been deployed on six VLA antennas. Observations at 4-band are also available through the RSRO program.

          Finally, the commensal VLITE system will take data at P-band during regular observations that use bands other than P-band.  The VLITE system is deployed on ten VLA antennas.  Observers wishing to gain access to the commensal VLITE data taken during their VLA observations should follow the instructions on the VLITE web page for doing so.


          Documentation Summary

          https://science.nrao.edu/observing/call-for-proposals/2016A - The July 1, 2015 NRAO call for proposals for the August 3, 2015 deadline (2016A).

          https://my.nrao.edu - The NRAO User Portal. This is a gateway to the NRAO interactive services that include the Proposal Submission Tool (PST).

          https://science.nrao.edu/facilities/vla/docs/manuals/proposal-guide/pst - The NRAO Proposal Submission Tool (PST) online manual.

          https://obs.vla.nrao.edu/ect - The VLA Exposure Calculator Tool (ECT).

          https://science.nrao.edu/facilities/vla/docs/manuals/propvla/gost - The VLA GO Setup Tool (GOST).

          https://science.nrao.edu/facilities/vla/docs/manuals/obsguide - The VLA Observing Guide.


          To report errors or problems encountered in any link or while using any NRAO tool listed here, please submit a ticket through the NRAO Helpdesk.

           

           

          Many thanks to all the VLA staff and our RSRO participants who have worked long and hard to commission these capabilities and who have helped to create this extensively updated set of documentation.

          Performance of the VLA during the Next Semester

          Introduction

          This section contains details of the VLA's resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of radio interferometry and synthesis imaging are found elsewhere. In particular, see References 1 and 2, listed in the Documentation.

          Resolution

          VLA Resolution

          The VLA's resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is "blind" to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the VLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/Bmax) are not seen – the resulting image will be smoothed to the resolution of the array. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can partly or fully recover this missing information, which can only be obtained by observing in a more compact array configuration, by using the mosaicing method, or by utilizing data from an instrument such as a large single antenna or a compact array comprising smaller antennas which provides this information.

          Table 2 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power) for the central frequency for each band, and the scale at which severe attenuation of large scale structure occurs.

          These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith.
          Footnotes:
          1. Bmax is the maximum antenna separation, Bmin is the minimum antenna separation, θHPBW is the synthesized beam width (FWHM), and θLAS is the largest scale structure "visible" to the array.
          2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (DnC, CnB, BnA) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., DnB, or CnA) is provided.
          3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θHPBW. The values for snapshots are about 1.3 times the listed values.
          4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two.
          5. For the C configuration an antenna from the middle of the north arm is moved to the central pad "N1". This results in improved imaging for extended objects, but may slightly degrade snapshot performance. Note that although the minimum spacing is the same as in D configuration, the surface brightness sensitivity and image fidelity to extended structure is considerably inferior to that of the D configuration.

          The following figure is a graphical representation of the synthesized beamwidths for natural and robust weighting for the four main array configurations between 1 and 50 GHz. Also available are synthesized beamwidth figures for the low frequency (1–12 GHz) and the high frequency (12–50 GHz) receiver bands.

          A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; the goal of implementing a new digital Pie Town link, now that EVLA construction is complete, remains unfunded.

          Sensitivity

          The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by:

           

           

          where:

          - SEFD is the "system equivalent flux density" (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEFD indicate more sensitive performance. For the VLA's 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62TsysA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band.
          - ηc is the correlator efficiency (~0.93 with the use of the 8-bit samplers).
          - npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in 'RCP' or 'LCP'.
          - N is the number of antennas.
          - tint is the total on-source integration time in seconds.
          - Δν is the bandwidth in Hz.

          Figure 1 shows the SEFDs as a function of frequency used in the VLA exposure calculator for those Cassegrain bands currently installed on VLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 3 gives the SEFD at some fiducial VLA frequencies.

          Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical:

          • For the more commonly-used "robust" weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values.
          • Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 3 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 3; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations.
          • Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band.
          • Losses from the 3-bit samplers.  The VLA's 3-bit samplers incur an additional 10 to 15% loss in sensitivity above that expected -- i.e., the efficiency factor ηc = 0.78 to 0.83.

           

            In general, the zenith atmospheric opacity to microwave radiation is very low - typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable.
            Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz.
            Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. However, the antenna tip capability is currently unavailable for the upgraded VLA -- it is hoped that this will again be available in the next year.
            When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item.
            • Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target.

              Both CASA and AIPS allow the application of elevation-dependent gains and an estimated opacity generated from ground-based weather (e.g., through the tasks gencal and plotweather in CASA, and INDXR in AIPS).

            • Pointing. The SEFD quoted above assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime can be much worse -- occasionally exceeding 1 arcminte, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz.
            To achieve better pointing, "referenced pointing" is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 - 3 arcseconds (but more typically 5 to 7 arcseconds) if the reference source is within about 15 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup.
            Use of referenced pointing is highly recommended for all Ku, K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.3 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution.

            To aid VLA proposers there is an on-line guide to the exposure calculator; the exposure calculator provides a graphical user interface to these equations.

            Special caveats apply for P-band (230 – 470 MHz) observing.  The listed SEFD in table 3 is from an observation taken far from the Galactic plane, where the sky brightness is about 30K.  At this band, Galactic synchrotron emission is very bright in directions near the Galactic plane.  The system temperature increase due to Galactic emission will degrade sensitivity by factors of two to three for observations in the plane, and by a factor of 5 or more at or near the Galactic center.   In addition, the antenna efficiency (currently about 0.31 for 300 MHz) will decline with both increasing and decreasing frequencies from the center of P-band.

            The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by:

            where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/Bmax), and S in mJy per beam, the parameter F depends on the synthesized beam, therefore on the array configuration, and has the approximate value F = 190, 18, 1.7, 0.16 for A, B, C, and D configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ΔIm, for S. Note that Equation 2 is a beam-averaged surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Documentation.

            For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791):

            where D is the distance to the galaxy in Mpc, and SΔV is the HI line area in units of Jy km/s.

            VLA Frequency Bands and Tunability

            Bands

            For observations taken with the 8-bit samplers,  each receiver can tune to two different frequencies, each 1.024 GHz wide,  within the same frequency band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies, except for the low-band receiver (50 – 500 MHz), which provides linear polarization (X and Y).  Each of these four data streams follows the VLA nomenclature, and are known as IF (for "Intermediate Frequency" channel) "A", "B", "C", and "D". IFs A and B provide RCP (or Y when applicable), IFs C and D provide LCP (or X when applicable).  IFs A and C are always at the same frequency, as are IFs B and D (but note that the A and C IFs  frequency is usually different from the B and D frequency). We normally refer to these two independent data streams as "IF pairs" – i.e., the 'A/C' pair and the 'B/D' pair.   In 8-bit mode, a maximum of 1.024 GHz can be correlated for each IF pair (see the WIDAR Section), for a total maximum bandwidth of  2.048 GHz.  To distinguish this 8-bit system from the 3-bit system, these IF pairs are denoted A0/C0 and B0/D0.

            With the 3-bit samplers, more options are available.  This system provides four (R,L) polarization pairs, each 2048 MHz wide.  The A/C IF pair provides two sampled pairs, labelled A1/C1 and A2/C2, and the B/D IF pair provides two sampled pairs, labelled B1/D1 and B2/D2.

            For more details on the 8-bit and 3-bit samplers see the VLA Samplers section.

            The tuning ranges, along with default frequencies for continuum applications, are given in Table 4 below.

            Table 4: Default frequencies for "continuum" applications
            BandRange18-bit continuum applications (GHz)3-bit continuum applications (GHz)

            (GHz)IF pair A0/C0IF pair B0/D0IF pair A1/C1IF pair A2/C2IF pair B1/D1IF pair B2/D2
            4 m (4) 0.058 – 0.0842 .054 – .086
            90 cm (P) 0.23 – 0.473 0.224 – 0.4803
            20 cm (L) 1.0 – 2.04 1.0 – 1.54 1.5 – 2.04
            13 cm (S) 2.0 – 4.0 2.0 – 3.0 3.0 – 4.0
            6 cm (C) 4.0 – 8.0 4.5 – 5.5 5.5 – 6.5 4.0 – 6.0 6.0 – 8.0
            3 cm (X) 8.0 – 12.0 8.0 – 9.0 9.0 – 10.0 8.0 – 10.0 10.0 – 12.0
            2 cm (Ku) 12.0 – 18.0 13.0 – 14.0 14.0 – 15.0 12.0 – 14.0 14.0 – 16.0 16.0 – 18.0
            1.3 cm (K) 18.0 – 26.5 20.2 – 21.2 21.2 – 22.2 22.0 – 24.0 24.0 – 26.0 18.0 – 20.0 20.0 – 22.0
            1 cm (Ka) 26.5 – 40.0 32.0 – 33.0 31.0 – 32.0 33.0 – 35.0 35.0 – 37.0 29.0 – 31.0 31.0 – 33.0
            0.7 cm (Q) 40.0 – 50.0 40.0 -- 41.0 41.0 – 42.0 44.0 – 46.0 46.0 – 48.0 40.0 – 42.0 42.0 – 44.0
            Notes:
            1.  Listed here are the nominal band edges.  For all bands, the receivers can be tuned to frequencies outside this range, but at the cost of diminished performance.  Contact VLA staff for further information.
            2. The 4-band system is currently under development. Observing time may be requested through the RSRO program.
            3. The default setup for P-band will provide  16 subbands from the A0/C0 IF pair, each 16 MHz wide, to cover the frequency range 224 – 480 MHz. The channel resolution is 125 kHz. 
            4. The default frequency set-up for L-band comprises two 512 MHz IF pairs (each comprising 8 contiguous subbands of 64 MHz) to cover the entire 1 – 2 GHz of the L-band receiver.

            Tuning Restrictions

            In general, for all frequency bands except Ka, if the total span of the two independent IF pairs of the 8-bit system (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For K, Ka and Q bands (the only bands where a span greater than 8 GHz is possible), there are special rules:

            • At Ka band, the low frequency edge of the A0/C0 IF pair must be greater than 32.0 GHz. There is no restriction on the B0/D0 frequency, unless the B0/D0 band overlaps the A0/C0 band when the latter is tuned at or near the 32.0 GHz limit.  In this case, the Observation Preparation Tool (OPT) may not allow the requested frequency setups.  Users wanting to use such a frequency setup are encouraged to contact the NRAO Helpdesk for possible tuning options.
            • At K and Q bands, if the frequency span is greater than 8.0 GHz, the B0/D0 frequency must be lower than the A0/C0 frequency.

            For the 3-bit system the maximum frequency span permitted for the A1/C1 and A2/C2 IF pairs is about 5000 MHz.  The same restriction applies to B1/D1 and B2/D2.  The tuning restrictions given above for the separation and location of the 8-bit pairs A0/C0 and B0/D0 also apply to the 3-bit pairs, with A0/C0 replaced by A1/C1 and A2/C2, and B0/D0 replaced by B1/D1 and B2/D2.

            VLA Samplers

            The VLA is equipped with two different types of samplers, 8-bit, 1GHz bandwidth, and 3-bit, 2GHz bandwidth. The choice depends on your science goals, and on technicalities described below.

            The 8-bit Set consists of four 8-bit samplers running at 2.048 GSamp/sec.  The four samplers are arranged in two pairs, each pair providing 1024 MHz bandwidth in both polarizations.  The two pairs are denoted A0/C0 and B0/D0.  Taken together, the four samplers offer a maximum of 2048 MHz coverage with full polarization. The frequency spans sampled by the two pairs need not be adjacent.  Some restrictions apply, depending on band, as described in the section on Frequency Bands and Tunability.

            The 3-bit Set consists of eight 3-bit samplers running at 4.096 GSamp/sec.  The eight samplers are arranged as four pairs, each pair providing 2048 MHz bandwidth in both polarizations.  Two of these pairs, denoted A1/C1 and A2/C2 cannot span more than 5000 MHz (lower edge of one to the higher edge of the other). The same limitation applies to the second pair, denoted B1/D1 and B2/D2.  The tuning restrictions are described in the section on  Frequency Bands and Tunability.

             

            Which set to use?

            • S, L, and 4/P-band observations, whether line or continuum, should use the 8-bit sampler set.
            • C and X-band continuum observations should use 3-bit samplers in order to exploit the full 4 GHz bandwidth: in spite of the 15% reduction in sensitivity that comes with 3-bit (at equal bandwidth to the 8-bit samplers; see below for details) and the reduced effective bandwidth after removing RFI this still provides superior overall sensitivity. For more details we refer to EVLA memo 166.
            • Ku, K, Ka, and Q band continuum observations should use the 3-bit samplers for maximum bandwidth.
            • Wide-band spectral line searches requiring more than 2 GHz span should use the 3-bit samplers.
            • Spectral-line observations which fit within two (possibly disjoint) 1 GHz bands should use the 8-bit set.
            • Simultaneous continuum and high resolution spectral line observation can use mixed 3-bit and 8-bit samplers. The 3 bit samplers in this case will be set up to deliver the continuum data, while the 8-bit samplers will be for the spectral line data. The mix mode can be used in C and the higher frequency bands.


            Major Characteristics of each Set

            The 8-bit samplers are warranted for observations at 4/P-band, L-band, and S-band. The full analog bandwidth from the receivers fits within the 2048 MHz span covered by the samplers.

            For the 3-bit samplers users need to be aware of the following issues:

            • Sensitivity: compared to the 8-bit system, the sensitivity of  the 3-bit samplers is worse by ~15% (at equal bandwidth). Alternatively, a given continuum noise level, requiring on-source  integration time T with the 8-bit (two bands of 1GHz), requires  0.33T with the 3-bit (4 bands of 2GHz, assuming the bandwidth is available from the front end.)
            • Resonances: each of the eight 3-bit samplers on an antenna has a resonance about 3 MHz wide.  Each resonance is independent of all others, so there is no correlated signal between antennas.  The resonance degrades the spectrum in its narrow frequency range, but has little effect on continuum observing.  Bandpass solutions will be affected, but can be interpolated over. Spectral-line calibration and  images at  the affected frequencies will show significant loss in sensitivity. The resonances are easily seen in autocorrelation spectra, and it is recommended that users, especially spectral-line users, utilize these to locate the compromised frequencies.
            • Amplitude Calibration: The traditional method for both 8- and 3-bit systems is to observe a flux-density calibrator, use self-cal to determine the antenna amplitude calibration factors (gains), and transfer the gains to the phase calibrator and target. For 3-bit samplers this procedure gives results good to  5%, between elevations of 20-70degrees. (Expect worse at the upper edge of Q-band and/or during bad weather). The switched power data can be used to correct for system gain variations and works well for the 8-bit samplers.  For 3-bit samplers, the Pdif depends on the Psum,  i.e. Pdif is non-linear, and its application will bias the resulting visibilities by 5 to 10%.  The origin of this effect is understood, but we have not yet determined how best to compensate for it.  Because of this, we do not recommend use of the 'Psum' and 'Pdif' data to calibrate visibilities from the 3-bit samplers. We do, however, recommend that the 'Requantizer Gains' in the switched power data be applied to remove gain changes. For more information about the switched power, and 'Psum' and 'Pdif, see EVLA memo 145.

             

            Setting up the 8-bit or 3-bit Samplers

            Either set requires an initial scan for each individual LO (frequency) tuning, during which power levels are optimized.

            For the 8-bit system, a 'dummy scan' of 1 minute duration is sufficient for each tuning.  This  is usually done while the antennas are slewing at the start of an observing file, as the pointing direction of the antennas is not critical.

            For the 3-bit system, the requirements are more demanding, see the section on 3-bit set-up. The minimum setup time is 1 minute for each tuning, to adjust the power levels and bandpass slopes across the 2GHz samplers.  These values are retained and applied if the tuning is re-encountered in the same observation.  In addition, every time the LO setup is changed, whether or not it is new (e.g. changing from 8-bit X-band reference pointing back to target),  a scan of 30 seconds is needed to re-set the subband gains (requantizers) in the correlator.  For better amplitude calibration at high frequencies, the 3-bit initial setup should be near the elevation of the target, so do it after the first 8-bit setup described above. For 3-bit observing without 8-bit (e.g. C or X-band, without reference pointing), the power variation with elevation is small so the 3-bit setup can be done at any elevation.

            For settings that use a mix of 3-bit and 8-bit samplers, the guidelines to set up the 3-bit samplers should be followed.


            Other issues

            The overhead for setup of 3-bit samplers can eat into observing time, especially for projects with many different LO settings, and/or sources all over the sky, accompanied by a band change, reference pointing, and requantizer reset for each direction. The impact is most severe for short scheduling blocks.

            Polarization testing conducted so far indicates no degradation of performance by using the 3-bit samplers.

             

            Acknowledgements

            NRAO is grateful to Professor Rob Ivison for supporting the upgrade of some of the 3-bit samplers on the VLA via a grant from the European Research Council.  For observations using the 3-bit samplers between May 2015 and March 2018 we encourage users to include the following text in the Acknowledgments section of their publications:

            "We acknowledge funding towards the 3-bit samplers used in this work from ERC Advanced Grant 321302, COSMICISM."

            Field of View

            Primary Beam

            The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θPB = 45/νGHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as "mosaicing," in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 in Documentation for details.

            Guidelines for mosaicing with the VLA are given in the Guide to Observing with the VLA

            Chromatic Aberration (Bandwidth Smearing)

            The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array.

            These effects can be parameterized by the product of the fractional bandwidth (Δν/ν0) with the source offset in synthesized beamwidths (θ0HPBW). Table 5 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 5 should be used to determine how much spectral averaging can be tolerated when imaging a particular field.

            Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). Δν/ν0 is the fractional bandwidth; θ0HPBW is the source offset from the phase tracking center in units of the synthesized beam.

            Time-Averaging Loss

            The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parametrized, except for observations at δ = 90°. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with Δν/ν0 replaced by ωeΔtint, where ωe is the Earth's angular rotation rate, and Δtint is the averaging interval.

            For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 (in Documentation) considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 6, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θΔtint)2, where θ is the distance from the phase center and Δtint is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see the section on Time Resolution and Data Rates) in the A and B configurations.

            Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines.

            Non-Coplanar Baselines

            The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations.

            Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the "W-projection" method available in CASA; see EVLA Memo 67 for more details.

            Time Resolution and Data Rates

            The default integration times for the various array configurations are as follows:

            ConfigurationObserving
            Bands
            Default
            integration time
            D, C L S C 5 seconds
            D, C X Ku K Ka Q 3 seconds
            B all 3 seconds
            A all 2 seconds

            Observations with the 3-bit (wideband) samplers must use these integration times. Observations with the 8-bit samplers may use shorter integration times, but these must be requested and justified explicitly in the proposal, and obey the following restrictions:

            Minimum Integration Times and Maximum Data Rates
            Proposal type

            Minimum integration time

            Maximum data rate
            General Observing (GO) and
            Shared Risk Observing (SRO)
            50 msec 25 MB/s (90 GB/hr), or up to
            60 MB/s (216 GB/hr) with additional justification
            Resident Shared Risk Observing (RSRO) < 50 msec > 60 MB/s (216 GB/hr)

            Note that integration times as short as 5 msec and data rates as high as 300 MB/s can be supported for some observing, though any such observing is considered Resident Shared Risk Observing (RSRO). For these short integration times and high data rates there will be limits on bandwidth and/or number of antennas involved in the observation. Those desiring to utilize such short integration times and high data rates should consult with NRAO staff.

            The maximum recommended integration time for any EVLA observing is 60 seconds. For high frequency observations with short scans (e.g., fast switching, as described in Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)), shorter integration times may be preferable.

            Observers should bear in mind the data rate of the VLA when planning their observations. For Nant antennas and integration time Δt, the data rate is:

            Data rate ~ 45 MB/sec × (Nchpol/16384) x Nant × (Nant - 1)/(27×26) / (Δt/1 sec)
            ~ 160 GB/hr × (Nchpol/16384) x Nant × (Nant - 1)/(27×26) / (Δt/1 sec)
            ~ 3.7 TB/day × (Nchpol/16384) x Nant × (Nant - 1)/(27×26) / (Δt/1 sec)
            ...equation (4)

            Here Nchpol is the sum over all subbands of spectral channels times polarization products:

            Nchpol = Σsb Nchan,i x Npolprod,i

            where Nchan,i is the number of spectral channels in subband i, and Npolprod,i is the number of polarization products for subband i (1 for single polarization [RR or LL], 2 for dual polarization [RR & LL], 4 for full polarization products [RR, RL, LR, LL]). This formula, combined with the maximum data rates given above, imply that observations using the maximum number of channels currently available (16384) will be limited to minimum integration times of ~2 seconds for standard observations, and 0.8 seconds for shared risk observations.

            These data rates are challenging for data transfer, as well as data analysis. Currently data may either be downloaded via ftp within the Science Operations Centers, or mailed on hard drives for those not in the same building as the archive. The Archive Access Tool allows some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations.

            Higher time resolutions and data rates are possible in principle but will be considered only through the Resident Shared Risk program.

            Note: The data rate formula given above does not account for the auto-correlations delivered by WIDAR. Precise data rate values can be obtained through the use of the General Observing Setup Tool (GOST), or the Resource Catalog of the Observation Preparation Tool.

            Radio-Frequency Interference

            The very wide bandwidths of the upgraded Very Large Array mean that RFI (radio-frequency interference) will be present in a far larger fraction of VLA observations than in observations made with the old systems.  Considerable effort has gone into making the VLA's new electronics as linear as possible, so that the effects of any RFI will remain limited to the actual frequencies at which the RFI exists.   Non-linear effects, such as receiver saturation, should occur only for those very unlikely, and usually very brief, times when the emitter is within the antenna primary beam.

            RFI is primarily a problem within the low frequency bands (C, S, L, and the low-band system), and is most serious to the D configuration.  With increasing frequency and increasing resolution comes an increasing fringe rate, which is often very effective in reducing interference to tolerable levels.

            The bands within the tuning range of the VLA which are allocated exclusively to radio astronomy are 1400-1427 MHz, 1660-1670 MHz, 2690-2700 MHz, 4990-5000 MHz, 10.68-10.7 GHz, 15.35-15.4 GHz, 22.21-22.5 GHz, 23.6-24.0 GHz, 31.3-31.8 GHz, and 42.5-43.5 GHz. No external interference should occur within these bands.

            RFI seen in VLA data can be internal or external.  Great effort has been expended to eliminate all internally-generated RFI.  Nevertheless, some internal RFI remains, which we are working hard to eliminate.   Nearly all such internally-generated signals are at multiples of 128 MHz.  So far as we know, all such internal signals are unresolved in frequency, and hence will affect only a single channel.

            Radio frequency interference of external origin will be an increasing problem to astronomical observations. Table 7 lists some of the sources of external RFI at the VLA site that might be observed within the VLA's expanded tuning range within L and S bands. Figure 3 shows a raw power cross-power spectrum at L-band. Figure 4 shows a similar plot for the lower half of S-band.

             

             

             

            The three last entries in the table deserve extra discussion.  These are all satellite transmissions, whose severity is a strong function of the angular offset between the particular satellite and the antenna.  It appears that significant degradation can occur if the antennas are within ~10 degrees of the satellite.  The great majority of the satellites are along the 'Clarke Belt' -- the zone of geosynchronous satellites.  As seen from the VLA, this belt is at a declination of about -5.5 degrees.  There are dozens -- probably hundreds -- of satellites 'parked' along this belt, transmitting in many bands:  S, C, Ku, K, and Ka at a minimum.  Observations of sources in the declination rate of +5 to -15 degrees can expect to be significantly degraded due to satellite transmission.  The Sirius digital radio system (and probably the satellites in the 2178 -- 2195 MHz band) comprises three satellites in a 24-hour, high eccentricity orbit with the apogee above the central U.S.  For the Sirius system, the orbit is arranged such that each of the three satellites spends about eight hours near an azimuth of 25 degrees and an elevation of 65 degrees.  The corresponding region in astronomical coordinates is between declinations 50 to 65 degrees, and hour angles between -1 and -2 hours.  Observations at S-band within that area may -- or may not -- be seriously affected.  The most reliable way to judge the seriousness of the satellite emissions is to inspect the switched power table data, particularly in those subbands where there is little RFI.

            VLA staff periodically observes the entire radio spectrum, with the VLA,  from 1.0 through 50.0 GHz with 125 kHz channel resolution to monitor the ever-changing RFI spectrum.  Plots from this program, accompanied with tables of identified sources are available at the RFI section of the VLA Observing Guide. Users concerned about the precise frequencies of strong RFI, and the likelihood of being affected, are encouraged to peruse these plots.

            Although most of the stronger sources of RFI are always present, it is very difficult to reliably predict their effect on observations.  Besides the already noted dependence on frequency and array configuration, there is another significant dependency on sky location for those satellites in geostationary orbit.  For these transmitters, (for example, the frequency range from 3.8 to 4.2 GHz), the effect on observing varies dramatically on the declination of the target source.  Sources near zero declination will be very strongly affected, while observations north of the zenith may well be nearly unaffected, especially at the highest resolutions.

            Also available are total-power plots of all RFI observations made by the interference protection group, from 1993 onwards at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu.

            The VLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it is possible to observe in spectral windows containing RFI, provided the spectra are well sampled to constrain Gibbs ringing, and spectral smoothing (such as Hanning) is applied.  Both AIPS and CASA provide useful tasks which automatically detect and flag spectral channels/times which contain strong RFI.

            Extracting astronomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish and subtract RFI signals from interferometer data is ongoing.

            The 3-bit samplers will be more susceptible to RFI signals than the 8-bit samplers, since the latter have more 'levels' within which these strong signals can be accommodated.  However, since the RFI power at the bands where the 3-bit samplers will most commonly be used (C, X, Ku, K, Ka, and Q) is nearly always less than the total noise power, we do not expect problems when wide-band 3-bit observing is done in these bands.

            Calibration of VLA data when strong RFI is present within a subband can be difficult.  Careful editing of the data, using newly available programs within CASA and AIPS, will be necessary before sensible calibration can be done.  The use of spectral smoothing (typically, Hanning), prior to editing and calibration, is strongly recommended when RFI is present within a subband.

            Identification and removal of RFI is always more effective when the spectral and temporal resolutions are high.  However, the cost of higher spectral and temporal resolution is in database size and, especially, in computing time.  A good strategy is to observe with high resolution, then average down in time and frequency once the editing is completed.

            Subarrays

            The continuum subarray option offers two 1 GHz baseband pairs with the 8-bit samplers in up to 3 subarrays, with the same spectral channel and polarization product options as are available for wideband observing. The setup for each subarray is completely independent, in terms of observing frequency, polarization products, and integration times.

            When using three subarrays, there are some restrictions on the number of antennas in each subarray. The Baseline Board in the correlator treats each set of 4 antennas independently, using a separate column of correlator chips. With 8 such columns, the correlator can handle up to 8x4= 32 antennas. The correlator configuration software requires that a given column not be split across subarrays. This does not matter when using only two subarrays, but forces some subtle restrictions when using three. For instance, one cannot observe with 9 antennas in each of 3 subarrays, because 9 antennas requires three columns (two with 4 antennas each, and one with 1 antenna); three subarrays of 9 antennas each would require 3x3= 9 columns, one more than are actually available. Splitting the array into 10, 9, and 8 antennas is allowed, since the first two subarrays use 3 columns each, while the third uses only two.

            The following table gives four examples of how correlator resources can be split into multiple subarrays. Antennas in each subarray are color-coded: red for subarray 1, green for subarray 2 (if present), and blue for subarray 3 (if present).  The last column gives the number of antennas in each subarray (e.g., in the setup shown in the first row, subarray 1 has 10 antennas; subarray 2 has 9 antennas; and subarray 3 has 8 antennas). In all cases a total of 27 antennas are used.  [The columns are numbered in reverse order (C7 to C0) to match the numbering scheme used on the actual Baseline Boards.]

            Some Possible Subarray Options
            Number of antennas correlated using each Baseline Board columnNumber
            C7 C6 C5 C4 C3 C2 C1 C0 of antennas
            4 4 2 4 4 1 4 4 10 + 9 + 8
            4 4 4 2 4 4 4 1 14 + 13
            4 4 3 4 4 3 3 2 11 + 11 + 5
            4 4 4 4 4 4 3 -- 27

            Positional Accuracy & Astrometry

            Summary: The position of a target can be determined to a small fraction of the synthesized beam, limited by atmospheric phase stability, the proximity of an astrometric calibrator, the calibrator-source cycle time, and the SNR on target.

            In preparation for observing, the a-priori position must be known to within the antenna primary beam, except perhaps for mosaicing observations. In the special case of using the phased VLA as a VLBI element, the a-priori position must be accurate to within the synthesized beam of the array.

            In post-processing, target positions are typically determined from an image made after phase calibration, i.e. correcting the antenna and atmospheric phases as determined on the reference source. The accuracy of the calibration determines the  accuracy of the positions in the image. (Note that phase self-calibration imposes the assumed position of the model, i.e.,  makes  the position indeterminate. Therefore, an absolute position cannot be determined after self-calibration, but relative positions between features within a self-calibrated image are valid.)

            It may help to think of astrometry in 2 steps, narrow and wide-field.

            In narrow-field astrometry, the target is close to the phase tracking center and the antennas nod every few minutes between the target and a calibrator. Under good conditions of phase stability, accurate antenna positions, (so-called 'baselines'), a strong target, a close calibrator with accurately known position, and rapid switching, the accuracy can approach 1-2% of the synthesized beam, with a floor of ~2 mas. Even under more typical conditions, 10% of the beam is readily achieved.

            Astrometric calibrators are marked 'J2000  A' in the VLA calibrator list, and have an accuracy of ~2 mas. Other catalogs from the USNO and the VLBA are also useful, but offsets may exist between the VLA and VLBA centroids, arising from extended structure in the particular source, and the different resolutions of the arrays.

            For studies of proper motion and parallax, the absolute accuracy of a calibrator may be less important than its stability over time. Close or in-beam calibrators with poor a-priori positions can be used, and tied to the ICRF reference frame in the same or separate observations.

            Phase stability can be assessed in real time from the Atmospheric Phase Interferometer (API) at the VLA site, which uses observations of a geostationary satellite at ~12GHz. Dynamic scheduling uses the API data to run a project under suitable conditions, specified by the user. Note that  VLBI projects using the phased VLA will typically be fixed date, not dynamically scheduled.

            The widefield case is to determine the positions of targets within the primary beam, referenced to a calibrator within the beam or close by. In addition to the previous effects, there are distortions as a function of position in the field, from small errors in the Earth orientation parameters (EOP) used at correlation time, differential aberration, and phase gradients across the primary beam.  With no special effort, the errors build up to roughly ~1 synthesized beam at a separation of ~10^4 beams from the phase tracking center.  Not all these errors are fully understood, and accurate recovery of positions over the full primary beam in the wideband, widefield  case is a research area. These effects are handled somewhat differently in the post-processing packages. Check with VLA staff for more details.

            Limitations on Imaging Performance

            Image Fidelity

            Image fidelity is a measure of the accuracy of the reconstructed sky brightness distribution. A related metric, dynamic range, is a measure of the degree to which imaging artifacts around strong sources are suppressed, which in turn implies a higher fidelity of the on-source reconstruction.

            With conventional external calibration methods, even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred.  The limiting factor is most often the effective phase stability of the telescope due to atmospheric/ionospheric fluctuations, although pointing errors and changes in atmospheric opacity can also be a limiting factor.  If a good model of the sky brightness distribution exist (e.g. use of compact structures of sufficient strength, though a good model of resolved sources in the field of view may also be used), standard self-calibration can be counted on to improve the images.  At low frequencies where the dominant phase error is due to ionospheric plasma density fluctuations, more advanced techniques may be required to account for change of ionospheric phase across the field of view.  Dynamic ranges in the thousands to hundreds of thousands can be achieved using these techniques, depending on the underlying nature of the errors. With the new WIDAR correlator and its much greater bandwidths and much higher sensitivities, self-calibration methods can be extended to observations of sources with much lower flux densities than very possible with the old VLA.

            The choice of image reconstruction algorithm also affects the correctness of the on-source brightness distribution. The CLEAN algorithm is most appropriate for predominantly point-source dominated fields. Extended structure is better reconstructed with multi-resolution and multi-scale algorithms. For high dynamic ranges with wide bandwidths, algorithms that model the sky spectrum as well as the average intensity can yield more accurate reconstructions.

            Invisible Structures

            An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Image reconstruction algorithms such as multi-resolution and multi-scale CLEAN can help to reduce or eliminate these negative bowls, but care must be taken in choosing appropriate scale sizes to work with.

            Table 2 gives the largest scale visible to each configuration/band combination.

            Poorly Sampled Fourier Plane

            Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in Documentation.

            Sidelobes from non-Deconvolved Sources

            At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task.  Sidelobe confusion is a strong function of observing band -- affecting most strongly L and P-band observations.  It is rarely a significant problem for observations at frequencies above 4 GHz.

            Sidelobes from Strong Sources

            An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere.

            Wide Field Imaging

            Wide-field observing refers primarily to the non-coplanar nature of the VLA when observing in non-snapshot mode. At high angular resolutions and low frequencies, standard imaging methods will produce artifacts around sources away from the phase center.  Faceted imaging (AIPS, CASA) and w-projection (CASA) techniques can be used to solve this problem.

            Another aspect of wide-field observing is the accurate representation of primary beam patterns, and their use during imaging. This is relevant only for very high dynamic ranges ( > 10000 ) or when there are very strong confusing sources at and beyond the half-power point of the primary beam.  This problem is worse with a wide-band instrument simply because the size of the primary beam (and the radius at which the half-power point occurs) varies with frequency, while there is also increased sensitivity out to a wider field of view. Work is under way to commission algorithms that deal with these effects by modeling and correcting for frequency-dependent and rotating primary beams per antenna, during imaging. Please note, however, that most advanced methods will lead to a significant increase in processing time, and may not always be required. Therefore, in the interest of practicality, they should be used only if there is evidence of artifacts without these methods.

            Finally, all of the above effects come into play for mosaicing, another form of wide-field imaging in which data from multiple pointings are combined during or after imaging.

            Wide-band Imaging

            The very wide bandpasses provided by the upgraded Very Large Array enable imaging over 2:1 bandwidth ratios -- at L, S, and C bands, the upper frequency is twice that of the lower frequency.   It is this wide bandwidth which enables sub-microJy sensitivity.

            In many cases, where the observation goal is a simple detection, and there are no strong sources near to the region of interest, standard imaging methods that combine the data from all frequencies into one single image (multi-frequency-synthesis) may suffice.  This is because the wide-band system makes a much better synthesized beam -- especially for longer integrations -- than the old single-frequency beam, thus considerably reducing the region of sky which is affected by incorrect imaging/deconvolution.  A rough rule of thumb is that -- provided a strong source is not adjacent to the target zone -- if the necessary dynamic range in the image is less than 1000:1, (i.e., the strongest source in the beam is less than 1000 times higher than the noise), a simple wide-band map may suffice.  

            For higher dynamic ranges, complications arise from the fact that the brightness in the field of view dramatically changes as a function of frequency, both due to differing structures in the actual sources in the field of view, and due to the attenuation of the sources by the primary beam.  One symptom of such problems is the appearance of radial spokes around bright sources, visible above the noise floor, when imaged as described above.  

            The simplest solution is to simply make a number of maps (say, one for each subband), which can then be suitably combined after correction for the primary beam shape. But with up to 64 subbands available with the VLA's new correlator, this is not always the optimal approach.  Further, images at all bands must be smoothed to the angular resolution at the lowest frequency before any spectral information can be extracted, and with a 2:1 bandwidth the difference in angular resolution across the band will be significant.  

            A better approach is to process all subbands simultaneously, utilizing software which takes into account the possibility of spatially variant spectral index and curvature, and knows the instrumentally-imposed attenuation due to the primary beam. Such wideband imaging algorithms are now available within CASA as part of the clean task, and work is under way to integrate them fully with wide-field imaging techniques.

            Calibration and Flux Density Scale

            The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators, and is also available within the Observation Preparation Tool.

            Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration - consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, S, C, X, Ku, K, Ka, and Q bands.

            Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the VLA cannot accurately measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 8 shows the flux densities of these sources in January 2012 at the standard VLA bands.  The flux density scale for the VLA, from 1 through 50 GHz, is based on emission models of the planet Mars, which is then calibrated to the CMB dipole using WMAP (Wilkinson Microwave Anisotropy Probe) observations (see Perley and Butler, 2013, for details).  The source 3C286 (=J1331+3030) is known to be non-variable, and has thus been adopted as the prime flux density calibrator source for the VLA. The adopted polynomial expression for the spectral flux density for 3C286 is:

            \[\log(S) = 1.2515 - 0.4605 \log(f) - 0.1715 \log^2(f) + 0.0336 \log^3(f)\]

            where S is the flux density in Jy, and f is the frequency in GHz.

            The absolute accuracy of our flux density scale is estimated to be about 2%.   With care, the internal accuracy in flux density bootstrapping is better than 1% at all bands except Q-band, where pointing errors limit bootstrap accuracy to perhaps 3%.   Note that such high internal accuracies are only possible in long-duration observations where the antenna gains curves and atmospheric opacity can be directly measured, and where there is good elevation overlap between the target source(s) and the flux density standard calibrator.

            Table 8: Flux densities (Jy) of Standard Calibrators for January 2012
            Source

            1465 MHz

            2565 MHz
            4885 MHz 8435 MHz 14965 MHz 22460 MHz 36435 MHz43340 MHz
            3C48 = J0137+3309 15.56 9.80 5.39 3.14 1.77 1.19 0.73 0.63
            3C138 = J0521+1638 8.71 6.17 4.02 2.78 1.89 1.46 1.03 0.92
            3C147 = J0542+4951 21.85 13.75 7.59 4.49 2.59 1.77 1.10 0.94
            3C286 = J1331+3030 14.90 10.03 7.34 5.09 3.39 2.52 1.75 1.53
            3C295 = J1411+5212 22.15 12.95 6.41 3.34 1.62 0.957 0.507 0.403
            NGC7027 1.62 3.59 5.38 5.79 5.62 5.42 5.18 5.04

            The sources 3C48, 3C147, and 3C138 are all slowly variable.  VLA staff monitor these variations on timescale of a year or two, and suitable polynomial coefficients are determined for them which should allow accurate flux density bootstrapping. These coefficients are updated approximately every other year, and are used in the AIPS task SETJY and in the CASA task setjy.

            The VLA antennas have elevation-dependent gain variations which are important to account for at the four highest-frequency bands.  Gain curves are determined by VLA staff, and the necessary corrections are applied to the visibility data when these data are downloaded from the archive. In addition to this, atmospheric opacity will also cause an elevation-dependent gain which is particularly notable at these four highest frequency bands.   At the current time, we do not have an atmospheric opacity monitoring procedure, so users should utilize the appropriate tasks available in both AIPS and CASA to estimate and correct for the opacity using ground-based weather data.  Correction of these gain dependencies, plus regular calibration using a nearby phase calibrator, should enable good amplitude gain calibration for most users.  Note that extraordinary attenuation by clouds can only be (approximately) corrected for by regular observation of a nearby calibrator.

            A better procedure for removing elevation gain dependencies uses the AIPS task ELINT.  This task will generate a 2nd order polynomial gain correction utilizing your own calibrator observations.  This will remove both the antenna and opacity gain variations, and has the decided advantage of not utilizing opacity models or possibly outdated antenna gain curves.  Use of this procedure is only practical if your observations span a wide range in elevation.

            By far the most important gain variation effect is that due to pointing.  Daytime observations on sunny days can suffer pointing errors of up to one arcminute (primarily in elevation).  This effect can be largely removed by utilizing the 'referenced pointing' procedure.  This determines the pointing offset of a nearby calibrator, which is then applied to subsequent target source observations.  It is recommended that this local offset be determined at least hourly, utilizing an object within 15 degrees of the target source -- preferentially at an earlier HA.  Studies show that the maximum pointing error will be reduced to about 7 arcseconds, or better.  VLA staff continue to work on improving this essential methodology.

            The VLA's post-amplifiers are not temperature stabilized, and exhibit significant gain changes between night and day, particularly at the four highest frequency bands.  Changes as large as 30% have been seen between night and day in calm, clear conditions!  These gain changes (and others caused by possible changes in attenuator settings) are monitored and will be removed with excellent accuracy by application of the internal calibration signal, whose results are recorded in the switched power table (SY table, in AIPS).  These corrections are not applied by default -- users who wish to correct for these gain changes must utilize the appropriate tasks in AIPS or CASA.   For the most accurate flux density bootstrapping, this table must be applied to the visibility data before calibration.  Gain bootstrapping better than 1% can be accomplished for the 8-bit sampler system after application of the switched power data.  For the 3-bit system there is an additional complication, as the values of the switched power data are sensitive to the total power, as well as the system gain.  VLA staff are currently working on a methodology to remove the total power dependency.  Not applying the switched power data will reduce bootstrapping accuracy to perhaps 10%, and possibly worse, if the observation of the flux density calibrator is not close in time to the local phase/amplitude calibrator.

            Complex Gain Calibration

            General Guidelines for Gain Calibration

            Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is difficult to define simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below.

            • Under decent conditions (no thunderstorms or ionospheric storms) tropospheric effects dominate at frequencies higher than about 4 GHz, ionospheric effects dominate at frequencies lower than about 4 GHz.
            • Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the C and D configurations at L and S bands, and in the D configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes - a couple of times per hour is generally sufficient.
            • If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale.  The enhanced sensitivity of the VLA guarantees, for full-band continuum observations, that every field will have enough background sources to enable phase self-calibration at L and S bands.  At higher frequencies, the background sky is not sufficient, and only the flux of the target source itself will be available.
            • The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an "S" code in the calibrator database, and a more distant calibrator with a "P" code, the nearby calibrator is usually the better choice.  A detailed description of calibrator codes is available in the Key to the calibrator list.
            • In clear and calm conditions, most notably in the summer, phase stability often deteriorates dramatically after about 10AM, due to small-scale convective cells set up by solar heating.  Observers should consider a more rapid calibration cycle for observations between this time and a couple hours after sundown.
            • At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See Rapid Phase Calibration and the Atmospheric Phase Interferometer (API).
            • Use the figure below to estimate how much time is minimally needed for each gain calibrator scan.  For instance, a 1 Jy calibrator and 4 MHz total bandwidth requires at least 30 seconds on source

             

            Minimum time required on a gain calibrator scan as a function of bandwidth and calibrator flux for the rather extreme case of upper Q-band.  Durations derived from this plot will definitely be sufficient for all other bands.

            Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)

            For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used for very small source-calibrator separations. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the old VLA this was known as "fast-switching." For the upgraded VLA it is just a loop of source-calibrator scans with short scan length. This technique "stops" tropospheric phase variations at an "effective" baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest VLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Note, however, that for the compact D-configuration, and a typical wind velocity, this "effective" baseline length is the same as, or larger than, the longest baseline in the array, and it is not worth the increased overhead of short cycle times. Under these circumstances it is sufficient to calibrate every 5-10 minutes to track the instrumental changes. The fast switching technique will also not work in bad weather (such as rain showers, or when there are well-developed convection cells - most notably, thunderstorms). It is also important to specify correctly the required tropospheric phase stability as measured by the Atmospheric Phase Interferometer at observe time (see below).

            Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from References 9 and 10 in Documentation.  See also the High Frequency Observing guide for additional recommendations on observing at high frequencies.

            An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data are heavily used for the dynamic scheduling of the VLA.

            Characteristic seasonal averages are represented in Table 9 below:

            Note: day indicates sunrise to sunset values; night indicates sunset to sunrise values.

            Polarization

            For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used for this at frequencies of ~3 GHz and higher, or 3C147 at frequencies abover ~10 GHz.  The table below shows the measured fractional polarization and intrinsic angle for the linearly polarized emission for these four sources in December 2010.  Note that 3C138 is variable -- the polarization properties are known to be changing significantly over time, most notably at the higher frequencies.  See Perley and Butler (2013b) for details.

            More information on polarization calibration strategy can be found in the VLA Observing Guide.

             

            Freq.3C48Pol3C48Ang3C138Pol3C138Ang3C147Pol3C147Ang3C286Pol3C286Ang
            GHz % Deg. % Deg. % Deg. % Deg.
            1.05 0.3 25 5.6 -14 <.05 - 8.6 33
            1.45 0.5 140 7.5 -11 <.05 - 9.5 33
            1.64 0.7 -5 8.4 -10 <.04 - 9.9 33
            1.95 0.9 -150 9.0 -10 <.04 - 10.1 33
            2.45 1.4 -120 10.4 -9 <.05 - 10.5 33
            2.95 2.0 -100 10.7 -10 <.05 - 10.8 33
            3.25 2.5 -92 10.0 -10 <.05 - 10.9

            33

            3.75 3.2 -84 - - <.04 - 11.1 33
            4.50 3.8 -75 10.0 -11 0.1 -100 11.3 33
            5.00 4.2 -72 10.4 -11 0.3 0 11.4 33
            6.50 5.2 -68 9.8 -12 0.3 -65 11.6 33
            7.25 5.2 -67 10.0 -12 0.6 -39 11.7 33
            8.10 5.3 -64 10.4 -10 0.7 -24 11.9 34
            8.80 5.4 -62 10.1 -8 0.8 -11 11.9 34
            12.8 6.0 -62 8.4 -7 2.2 43 11.9 34
            13.7 6.1 -62 7.9 -7 2.4 48 11.9 34
            14.6 6.4 -63 7.7 -8 2.7 53 12.1 34
            15.5 6.4 -64 7.4 -9 2.9 59 12.2 34
            18.1 6.9 -66 6.7 -12 3.4 67 12.5 34
            19.0 7.1 -67 6.5 -13 3.5 68 12.5 35
            22.4 7.7 -70 6.7 -16 3.8 75 12.6 35
            23.3 7.8 -70 6.6 -17 3.8 76 12.6 35
            36.5 7.4 -77 6.6 -24 4.4 85 13.1 36
            43.5 7.5 -85 6.5 -27 5.2 86 13.2 36

            High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA's new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible.

            The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing.

            Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive - 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. Ionosphere correction can also be performed in CASA using the task gencal; consult the calibration chapter of CASA Cookbook for more details.

            VLBI Observations

            The VLA can participate in VLBI observations.  This is only allowed in phased array mode (single dish is only available through the VLBA Resident Shared Risk Observing program) with restricted WIDAR correlator configurations.  Also note that currently P-band cannot be phased.  For more details see the VLBI at the VLA documentation.  In phased array mode the program TelCal derives the antenna-based delay and phase corrections needed for antenna phasing in real time.  This correction is applied to the antenna signals before they are summed, requantized to 2-bits, and recorded in VDIF format on the Mark5C disk at the VLA site.  The disk(s) are then transported to Socorro, NM and correlated on the DiFX correlator with other VLBI stations which participated in the observation.  Standard VLA data, i.e., correlations between VLA antennas,  are also archived in the NRAO science data archive.

            Snapshots

            The two-dimensional geometry of the VLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest.

            Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm and longer wavelengths in the D configuration. For instance, at 20 cm a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. The deconvolution of the data is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing "Faint Images of the Radio Sky at Twenty-centimeters survey (FIRST, http://www.cv.nrao.edu/first/)" (B configuration), or the NRAO VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) (D configuration, all-sky) surveys.

            Shadowing and Cross-Talk

            Observations at low elevation in the C and D configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA tasks importasdm and importevla) when using the default inputs. AIPS task UVFLG can be used to flag VLA data based on shadowing, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. The CASA task flagdata can also be used to flag data based on shadowing.

            Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. This effect is important at low frequencies in compact configurations.  Careful examination of the visibilities is necessary to identify and remove this form of interference. The affected data would show time-variable high-amplitude points.

            Combining Configurations and Mosaicing

             

            Any single VLA configuration will allow accurate imaging of a range of spatial scales determined by the shortest and longest baselines. For extended and structured objects, it may be required to obtain observations in multiple array configurations. It is advisable that the frequencies used be the same for all configurations to be combined. The ideal combination of arrays results in a uv-plane with all cells equally filled by uv-points. To first order, this can be achieved by using the beam sizes of the individual arrays to inversely scale the on-source integration time. This approach is equivalent to achieving the same surface brightness sensitivity for all arrays on all scales. For the VLA, observations in the different configurations generate beam sizes that decrease by factors of 3, i.e. C configuration generates a 3 times smaller beam than D configuration, B 3 times smaller than C, and A three times smaller than B. Thus, on-source integrations would increase by about an order of magnitude between each array. Such a drastic increase is very expensive and, in fact, not necessary since some spatial scales are common to more than a single array, which is equivalent to some uv-cells being filled more than others. The best way to fill the uv plane depends on many factors, such as declination of the source, LST time of the observation, and bandwidth.

            For the VLA, experience shows that a factor of about 3 in on-source integration time for the different array configurations works well for most experiments.

            E.g., 20min on-source time in D, 1h in C, 3h in B, and 9h in A should produce a decent map. Using large bandwidths and multi-frequency synthesis will broaden all uv tracks radially and one may need even less array configurations or shorter integration times between the different arrays.

            Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing.  The VLA has no limit on the number of pointings for each mosaic. Typically hexagonal, rectangular, or individual pointing patterns are used and the overlap regions will result in an improved rms over each individual pointing. Given the many, potentially short observations, it is important to obey the data rate limits outlined in the Time Resolution and Data Rates Section. On-the-fly (OTF) mosaics, i.e. dumping the data while moving the telescopes across the source, is also available.

            Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 (Documentation) for more information.

            Guidelines for mosaicing with the VLA are given in the VLA Observing Guide.

            Pulsar Observing

            The VLA can be used for two kinds of pulsar observing: phase-binning using the WIDAR correlator, and using the phased-array for pulsar dedispersion (either search or fold mode). Either of these types of observing are considered Resident Shared Risk (RSRO), and close collaboration with NRAO staff is required for their use.

            The WIDAR Correlator

            Introduction

            The correlator configurations offered for general observing may be divided into three basic modes: wideband, spectral line, and subarrays. Note that the possible setups are also subject to the integration time and data rate restrictions outlined in the section on Time Resolution and Data Rates. The possibilities and restrictions are embodied in the General Observing Setup Tool (GOST) and in the Resources section of the Proposal Submission Tool (PST), which must be used to define the correlator configuration for General Observing (GO) and Shared Risk Observing (SRO) proposals.

            Note that phased array configurations are only allowed as part of VLBI experiments (see the section on VLBI Observations) or as Resident Shared Risk observations.

            Wideband Observing

            The wideband observing setups provide the widest possible bandwidth for a given observing band, with channel spacing depending on the number of polarization products as listed in the following table:

            Wideband & Subarray Correlator Options (all but P- and L-bands)

            Polarization products Channel spacing
            Full (RR, RL, LR, LL) 2   MHz
            Dual (RR, LL) 1   MHz
            Single (RR or LL) 0.5 MHz

            8-bit wideband setups are available for all observing bands, providing a total of 2 GHz of bandwidth per polarization (1 GHz per polarization at L-band, and 256 MHz per polarization at P-band).  3-bit setups are available for all bands above S-band, providing total bandwidths per polarization of 4 GHz (C/X bands), 6 GHz (Ku band), or 8 GHz (K/Ka/Q bands). In all cases but P- and L-band each of the subbands is 128 MHz wide. At L-band the default is 64 MHz/subband, yielding channels twice as narrow as those listed in the table above, while at P-band the default is 16 MHz/subband,  resulting in 125 kHz channel spacing.

            In many frequency bands the total processed bandwidth is less than that delivered by the front-end. In those cases the observer may independently tune two 1 GHz baseband pairs when using the 8-bit samplers, or four 2 GHz baseband pairs when using the 3-bit samplers, or choose to have a mix 8-bit and 3-bit samplers. The tuning restrictions are described in the section on VLA Frequency Bands and Tunability, and the 8-bit and 3-bit samplers are described in the section on VLA Samplers.

            Spectral Line Observing

            Basebands and Subbands

            Currently observers have access to very flexible correlator configurations using up to 64 subbands in up to 4 basebands sampled with the 8-bit and/or the 3-bit samplers.  These capabilities may be summarized as follows:

            • Two 1 GHz baseband pairs using the 8-bit samplers, or four 2 GHz baseband pairs using the 3-bit samplers, independently tunable within the limits outlined in the section on VLA Frequency Bands and Tunability. The 8-bit baseband pairs are referred to as A0/C0 and B0/D0, while the 3-bit samplers are A1/C1, A2/C2, B1/D1, and B2/D2.  The AC/BD nomenclature corresponds to that of the IF pairs in the pre-expansion VLA.
            • Up to 16 subband pairs (spectral windows) in each 3-bit baseband pair, and up to 32 subbands in each 8-bit baseband pair with a total of up to 64 subbands in any correlator configuration
              • Tuning, bandwidth, number of polarization products, and number of channels can be selected independently for each subband
              • All subbands must share the same integration time
              • No part of a subband can cross a 128 MHz boundary
              • Subband bandwidths can be 128, 64, 32, ..., 0.03125 MHz (128 / 2n, n=0, 1, ..., 12)
            • The sum over subbands of channels times polarization products is limited to 16,384 (without recirculation).
              • These may be spread flexibly over subbands and polarization products, in multiples of 64: 64, 128, 192, 256, 384, ..., 16384 cross-correlation products.
              • Recirculation may be used to increase the number of channels per subband for subbands narrower than 128 MHz, Baseline Board stacking may be used to increase the number of channels per subband for setups requiring less than 64 subbands.
              • Assigning many channels to a given subband may reduce the total bandwidth and/or the total number of subbands available.

            The remainder of this section discusses the various limitations in more detail, including some examples to show how they come up in practice.

            Subband tuning restrictions

            Each subband may be placed anywhere within a baseband, with the caveat that no subband may cross a 128 MHz boundary. Mathematically:

            νBB0 + n*128 MHz <= νsbLow <= νsbHigh <= νBB0 + (n+1)*128 MHz

            where:

            νBB0 the lower frequency edge of the baseband;
            n= 0, 1, ..., 7 (, ..., 15) (i.e., any integer between 0 and 7 for 8-bit, between 0 and 15 for 3-bit);
            νsbLow                                   
            the lower edge of the subband (i.e., the subband center frequency minus half the subband bandwidth);
            νsbHigh                                 
            the upper edge of the subband (i.e., the subband center frequency plus half the subband bandwidth).

            So for example, if the baseband were tuned to cover 10000-11024 MHz, one could place a 64 MHz subband to cover 10570-10634 MHz, but not to cover 10600-10664 MHz (because that would cross the 128 MHz boundary at 10640 MHz). Note in particular that the center of a baseband is a boundary and no line should be observed at the baseband center.

            The figure below illustrates these restrictions:

            Correlator configuration figure: bandpass8jul12.png

            The black curve shows the analog filter response for an 8-bit baseband covering 1024 MHz, starting at νBB0. The dashed blue vertical lines show the 128 MHz boundaries; no subband can cross those boundaries.  128 MHz subbands are thus constrained to cover a region between two of those boundaries, and no finer tuning is possible.  Narrower subbands, like the 64 MHz subband shown here in red, can be shifted around arbitrarily within one of the 128 MHz "slots", but cannot cross any of these boundaries. (The dotted vertical red lines show the boundaries of the 64 MHz subband, while the solid curve shows an illustrative line within the subband.)

            The analog filter shape defining the baseband rolls off severely at one edge of the baseband, so the 128 MHz "slot" at that edge has reduced sensitivity.  The baseband edge is at the lowest sky frequency in the baseband when using upper sideband, and at the highest sky frequency in the baseband when using lower sideband.

            Subband bandwidths & the digital filter response

            The bandwidth for each subband may be selected independently, and can be any of 128/2n MHz, for n= 0, 1, ..., 12: 128, 64, 32, 16, 8, 4, 2, or 1 MHz, or 500, 250, 125, 62.5, or 31.25 kHz.

            The usable portion of the subband is set by three effects. First, as discussed above, the analog filters which define the baseband are not perfect, leading to lower sensitivity in the 128 MHz near the baseband edge for the 8-bit samplers

            Second, because the digital filters are not infinitely sharp, the rejected sideband leaks in at both edges of the subband. This leads to additional (aliased) noise, with a factor ~2 increase in the noise at the subband edges, dropping to a few per cent within a few per cent of the subband edge. The precise filter shape and noise increase is a complex but predictable function of the subband bandwidth and the subband tuning.

            The third effect stems from the offset frequencies used for sideband rejection in the WIDAR correlator. The local oscillators at the individual antennas are tuned to slightly different frequencies, with those offsets taken out in the correlator. This means that each antenna observes a slightly different sky frequency, and thus some baselines will not give an interesting correlation near one edge of the subband. The maximum frequency shift is currently set to 32*f0, with the fundamental f0 being set to f0= max(25.6 kHz*sbBW/128 MHz, 100 Hz). Here sbBW is the smallest subband bandwidth within the baseband. For the wider subband bandwidths the maximum frequency shift corresponds to <1% of that bandwidth, but for narrower subbands the effect can be severe. For instance, a 31.25 kHz subband has f0= 100 Hz, and a maximum frequency shift of 3.2 kHz -- 10% of the subband may be lost on some baselines.

            Spectral channels and polarization products

            Each subband (without recirculation enabled) can have a different number of channels and polarization products, subject to two limitations:

            1. For the ithsubband, the number of spectral channels can be:
              • 64 nBlBP,i with full polarization products (RR,RL,LR,LL)
              • 128 nBlBP,i with dual polarization products (RR,LL)
              • 256 nBlBP,i with a single polarization product (RR or LL)
              Here nBlBP,i= 1, 2, 3, 4, 5, ..., 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.
            2. The sum over all subbands of nBlBP,i must be less than or equal to 64, the number of Baseline Board pairs in the correlator. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64x256= 16,384 (without recirculation).

            Baseline Boards are the boards in the WIDAR correlator where the actual cross-multiplications are done. There are 128 Baseline Boards arranged as 64 Baseline Board pairs (BlBPs).  The limitations given here correspond to the capabilities of the individual boards, and the finite number of boards the correlator has.

            Limitation #1 corresponds to the following table of the options for subband bandwidth and spectral resolution when using nBlBP Baseline Board pairs for a subband:

            Subband Bandwidth and Spectral Resolution Options (without recirculation)
            Subband bandwidth &
            total velocity coverage
            Full polarization products
            (RR, RL, LR, LL)
            64nBlBP spectral channels

            Channel spacing:
            Dual polarization products
            (RR, LL)
            128nBlBP spectral channels
            Channel spacing:
            Single polarization product
            (RR or LL)
            256nBlBP spectral channels

            Channel spacing:
            128 MHz 38400/νGHz km/s 2000/nBlBP kHz 600/nBlBPGHz km/s 1000/nBlBP kHz 300/nBlBPGHz km/s 500/nBlBP kHz 150/nBlBPGHz km/s
            64 19200 1000 / nBlBP 300 / nBlBP 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP
            32 9600 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP
            16 4800 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP
            8 2400 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP
            4 1200 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 /n BlBP
            2 600 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP
            1 300 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP
            0.5 150 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP
            0.25 75 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP
            0.125 37.5 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP
            0.0625 18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
            0.0325 9.375 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP 0.122 / nBlBP 0.037 / nBlBP
            Subband bandwidth and spectral resolution options. Note that the table entries refer to the spacing between spectral channels -- that spacing is before any frequency smoothing, so these channels are notindependent.
            • nBlBP is the number of Baseline Board Pairs assigned to the subband.
            • Each subband may have a different number of spectral channels and polarization products, and each may be tuned independently.
            • There can be at most 16 subbands per baseband, and nBlBP must be an integer: 1, 2, 3, 4, 5, ..., 64.
            • The sum of nBlBP over all subbands must be less than or equal to 64.
            • Use of more than one BlBP for a subband may further restrict the number of subbands available in one or more of the basebands; see text for details.

            Here are four examples of allowed general observing setups which use all 64 BlBPs to produce the maximum number of channels times polarization products:

            Baseband Subband

            Pol'n

            Products

            Spectral

            channels

            nBlBP
            Example 1 A0/C0 sb0 RR 16384 64
            Example 2 A0/C0 sb0 RR 8192 32
            A0/C0 sb1 RR, LL 1024 8
            A0/C0 sb2 RR, LL 512 4
            B0/D0 sb0 RR, LL 2048 16
            B0/D0 sb1 RR,RL,LR,LL 256 4
            Example 3 A0/C0 sb0 RR 8192 32
            A0/C0 sb1 LL 1024 4
            A0/C0 sb2 RR, LL 1024 8
            A0/C0 sb3 RR,RL,LR,LL 1024 16
            A0/C0 sb4 RR,RL,LR,LL 256 4
            Example 4 A0/C0 sb0-5 RR,RL,LR,LL 64 6 x 1
            A0/C0 sb6 RR, LL 3840 1 x 30
            A0/C0 sb7 RR 768 1 x 3
            A0/C0 sb8 RR,RL,LR,LL 192 1 x 3
            B0/D0 sb0-2 RR,RL,LR,LL 64 3 x 1
            B0/D0 sb3 LL 768 1 x 3
            B0/D0 sb4 RR, LL 2048 1 x 16

             

            Recirculation

            Recirculation is a term to describe the method of increasing the number of spectral channels in a subband using correlator software (as opposed to Baseline Board stacking which uses correlator hardware, see below). Currently recirculation is achieved by limiting the subband bandwidth and thus only available for subbands less than 128 MHz wide.  When limiting the bandwidth in a subband, the correlator software can be directed to use the remaining CPU cycles on a Baseline Board pair to obtain more lags (in factors of two), running the data through the board for a second, third, etc., time; hence Recirculation.

            At some time in the future an alternative method of recirculation, using the extra CPU cycles freed up by increasing the integration time, would be made available. Recirculation by limiting of the subband bandwidth to increase the number of channels (in factors of two) was used in the pre-expansion VLA correlator.

            Recirculation versus Baseline Board stacking

            When faced with the choice between Recirculation and Baseline Board stacking (see below) to increase the number of channels in a subband for subbands narrower than 128 MHz we recommend the former, which is supported in observatory software (GOST, OPT). For 128 MHz subbands Baseline Board Stacking should be utilized to increase the number of channels.

            The current implementation of Recirculation is that for each halving of the subband bandwidth the number of channels in the subband may be doubled without having to trade off the use of other subbands. Because recirculation is achieved by limiting the subband bandwidth, it is not supported for 128 MHz subbands, whereas for 64 MHz subbands only a factor 2 recirculation is supported, etc. The maximum recirculation factor for a subband is 128/(subband bandwidth in MHz), and of course also subject to other configuration restrictions such as data rate.

            The juggling between the requested number of channels, subband bandwidth and the available number of Baseline Board pairs is dependent on the science goals and not easily formulated in a standard answer. However, if subbands of less than 128 MHz are used, Recirculation becomes an option for setups that can also be achieved with Baseline Board stacking. In such cases we suggest to use Recirculation where possible, and within the General Observing or Shared Risk requirements.  This frees up unused Baseline Board pairs for other use; alternatively, one becomes less dependent on all Baseline Board pairs being in working order.

            Recirculation with factors 8 to 64 is designated Shared Risk, and recirculation with factors over 64 Resident Shared Risk. The latter choice may have severe implications for the sensitivity as visibility integration time is used as trade-off. Ask the NRAO Helpdesk for more details.

            Baseline Board stacking

            As opposed to recirculation, which increases the number of channels in a subband by exploiting otherwise unused CPU resources, Baseline Board Stacking adds more channels to a subband by adding correlator hardware resources, i.e. using up more Baseline Board pairs.  Using Baseline Board stacking may therefore limit the number of subbands available in one or more of the basebands.  Understanding how this works requires understanding some of the details of the correlator hardware. That understanding is built into the GOST, and observers may simply use that tool to find out whether their particular setup will, in fact, work. But the results can be confusing without some understanding of the hardware constraints from which they arise.  These hardware constraints are complex, and most observers will not need to understand these details.  The following section is for those few who are attempting complex line experiments, and who find the GOST or the RCT restricting the number of subbands and/or channels they can use in unexpected ways  Most observers can skip it.

            Baseline Board Stacking and Correlator Use

            First let us consider how the correlator hardware is organized. The cross-multiplications in the WIDAR correlator are spread across 64 Baseline Board (BlB) pairs, arranged into 4 quadrants of 16 BlB pairs each. Each baseband is connected directly to one of those quadrants. In the simplest mode, each of the 16 BlB pairs of a quadrant handles the correlations for one of the 16 subbands of the corresponding baseband. Four basebands and four quadrants are required, in order to handle the full 8 GHz of bandwidth per polarization provided by the 3-bit (wideband) samplers: that 8 GHz is split into four basebands of 2 GHz each, with each baseband fed into a different BlB quadrant. Each BlB pair in that quadrant handles a subband of maximum bandwidth 128 MHz, so 16 BlB pairs handles 16 subbands for a total of 16x128 MHz= 2048 MHz.

            A single BlB pair produces 256 cross-correlations per baseline for a single subband, which can be used for a single polarization product (e.g., RR with 256 spectral channels), or two (RR+LL with 128 spectral channels each), or four (RR,RL,LR,LL, with 64 spectral channels each).

             

            When using the 8-bit samplers, the total bandwidth is only 2 GHz per polarization, split into two basebands of 1 GHz each. The simplest continuum setup uses only two quadrants, since there are only two basebands; and only 8 subbands are required to span the 8x128 MHz= 1024 MHz of each baseband. Three-quarters of the correlator BlB hardware remain unused.

             

             

            The spectral line mode allows access to these `extra' correlator resources through Baseline Board stacking: using multiple BlB pairs to process the same subband and produce more cross-correlations for that subband. This is done using crossbar switches which make the data for a single subband available to several BlB pairs. Those BlB pairs can then be used to produce more spectral channels for that subband, with n BlB pairs producing 256*n cross-correlations per baseline. The limit on the total number of cross-correlations (16384) stems from the total number of BlB pairs (64): 64x256= 16384.

            Unfortunately completely flexible crossbar switches are expensive, and could not be implemented in the VLA's new correlator. This means that one cannot route a given subband to a randomly-chosen BlB pair. The routings which are possible, are as follows:

            1. A subband in a baseband can be routed to any BlB pair within the corresponding quadrant.
            2. Data coming into a given BlB pair in one quadrant, can be routed to the corresponding BlB pair in any other quadrant.

             

            Routing option #1 means that one could use all the BlB pairs within a quadrant to correlate a single subband, yielding 16x256= 4096 cross-correlations for that subband:

             

            Routing option #2 means that one could use the BlB pairs in all 4 quadrants to correlate a single subband. One simple case would use 4 BlB pairs to correlate each of the 16 subbands in a single baseband, yielding 4x256= 1024 cross-correlations for each of those subbands. Note that in this case, no BlB pairs are left to correlate any data from the second baseband.

            Using routing option #2 does come with a subtle cost: assigning a BlB pair in quadrant X to correlate a subband corresponding to quadrant Y, removes that BlB pair from use in the baseband corresponding to quadrant X...and hence also removes the corresponding subband in that baseband. So getting more channels for a subband in one baseband, may prevent the use of a subband in a different baseband. To take a simple example, consider an experiment where one wishes to observe a single line in dual polarization with 512 channels (requiring 4 BlB pairs), plus as much continuum bandwidth as possible. Naively one would say there are 16 subbands in each baseband; one is used for the spectral line, so that leaves 16+15= 31 subbands, and with the widest subband bandwidth (128 MHz) the total available continuum should be 31x128 MHz= 3968 MHz per polarization. Actually however there are only 15+15 subbands available, or 30x128 MHz= 3840 MHz per polarization, because the spectral line subband has "eaten" one BlB pair corresponding to the other baseband:

             

            If the same spectral line required twice as many channels, this result in the loss of two subbands in both of the basebands:

            In some cases one may want to use a different routing, to use up subbands in one baseband in preference to another. For instance, the same spectral line setup (2048 cross-correlations for a single spectral line subband, plus as much continuum as possible) could be set up to allow 13 "continuum" subbands in the A0/C0 baseband, and the full 16 "continuum" subbands in B0/D0:

            Understanding these confusing constraints can help observers set up the VLA more effectively to achieve their scientific goals. For instance, for a mixed line+continuum experiment it works best to use the resource tools to set up the baseband tunings and subband channelization for the most important lines first, then add the desired continuum, and then see what correlator resources remain for any lines of secondary interest.

            The above examples all use BlB pair stacking in powers of 2, but this is not required.  To give some idea of more complex possibilities, the following tables give two examples of other possible configurations. The RCT display shows how the Baseline Boards are used to process the individual subbands. The cyan boxes show the Baseline Boards used to process data from baseband A0/C0, while the yellow boxes show Baseline Boards used to process data from baseband B0/D0.

            Complex Configuration Example #1
            BasebandSubbandPol'n productsSpectral channelsnBlBP
            A0/C0 sb0 RR 10240 40
            A0/C0 sb1 LL 768 3
            A0/C0 sb2 RR,LL 2176 17
            B0/D0 sb0 RR 256 1
            B0/D0 sb1 RR,LL 384 3
            RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

             

             

             

            Complex Configuration Example #2
            BasebandSubbandPol'n productsSpectral channelsnBlBP
            A0/C0 sb0 RR 4352 17
            A0/C0 sb1 RR, LL 1152 9
            B0/D0 sb0 RR,RL,LR,LL 192 3
            B0/D0 sb1 RR, LL 4480 35
            RCT display: corr-cfg-fig:sro2_8bit_ac17+9_bd3+35

             

            Note that the individual subbands can have different bandwidths, and those bandwidths may be chosen completely independently of the number of spectral channels in each subband. So for instance a subband with bandwidth 2 MHz and 1152 spectral channels would have a channel separation of 2 MHz/1152= 1.736 kHz; but the observer could equally well choose a bandwidth of 64 MHz for that subband, leading  to a channel separation of 64 MHz/1152=  55.56 kHz.

            Use of the 3-bit samplers further extends the possibilities.  Here is one example:

             

            3-bit Complex Configuration Example #1
            BasebandSubbandPol'n productsSpectral channelsnBlBP
            A1/C1 sb0-8 RR, LL, RL, LR 9 x 64 9 x 1
            A1/C1 sb9 RR, LL 1 x 1152 1 x 9
            A1/C1 sb10 RR 1 x 1792 1 x 7
            A1/C1 sb11 RR, LL 1 x 384 1 x 3
            A2/C2 sb0-11 RR, LL, RL, LR 12 x 64 12 x 1
            A2/C2 sb12 LL 1 x 768 1 x 3
            B1/D1 sb0-3 RR, LL, RL, LR 4 x 64 4 x 1
            B1/D1 sb4 RR, LL, RL, LR 1 x 320 1 x 5
            B2/D2 sb0-6 RR, LL, RL, LR 7 x 64 7 x 1
            B2/D2 sb7 RR, LL 1 x 640 1 x 5
            RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

             

            Once again, the GOST implements all of these constraints, and is generally smart enough to figure out the routing scheme that works best for your particular request.

            Documentation

            Documentation for VLA data reduction, image making, observing preparation, etc., can be found in various manuals.   Current manuals are available on-line. Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the VLA user, not listed here, are available from our website.

            1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: "Synthesis Imaging in Radio Astronomy." This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. The lectures of the 2014 workshop are available at the 14th Synthesis Imaging Workshop web site.
            2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series.
            3. A GUIDE TO OBSERVING WITH THE VLA: Describes details of how to observe with the VLA once you have been allocated time on the VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide). Including, e.g., special observing modes such as:
              1. OBSERVING WITH THE 8-BIT (up to 2 GHz bandwidth) & 3-BIT (up to 8 GHz bandwidth) SAMPLER SYSTEMS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up);
              2. SPECTRAL LINE OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/line); 
              3. HIGH FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/hifreq);
              4. LOW FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/lofreq);
              5. VERY LOW FREQUENCY OBSERVING (< 500 MHz) (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/vlofreq);
              6. POLARIMETRY (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/pol);
              7. MOSAIC OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/mosaicking);
              8. RADIO FREQUENCY INTERFERENCE (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/rfi);
              9. MOVING OBJECTS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/moving);
              10. VLBI at the VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/vlbi).
            4. *CASA COOKBOOK: The CASA "Cookbook" for use of the package for data reduction of VLA (& ALMA) data is available, along with other documentation, from the CASA home page (http://casa.nrao.edu).  See http://casa.nrao.edu/docs/cookbook/.
            5. VLA CASA Guides: Tutorials and data reduction examples of VLA data in CASA (https://casaguides.nrao.edu/index.php/EVLA_Tutorials)
            6. *AIPS COOKBOOK: The Astronomical Image Processing System (AIPS) software is able to fully calibrate VLA data and do most imaging operations.   The exception is the wide-band (bandwidth synthesis) deconvolution which is being developed in CASA only.   ALMA data may also be reduced in AIPS although the package is not fully qualified to calibrate data from the ALMA linearly-polarized feeds.   The "Cookbook" description for calibration and imaging under the AIPS system can be found near all public workstations in the SOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See http://www.aips.nrao.edu/cook.html.
            7. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See http://www.aips.nrao.edu/goaips.html.
            8. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibrators in both 1950 and J2000 epoch and a discussion of gain and phase calibration, and polarization calibration.  See https://science.nrao.edu/facilities/vla/docs/manuals/cal.
            9. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983.
            10. *HISTORICAL VLA MEMO SERIES: archive memo series from the early days of the VLA.  See http://library.nrao.edu/vlam.shtml.
            11. *RECENT VLA MEMO SERIES: the memo series relating to the expanded capabilities of the VLA.   See http://library.nrao.edu/evla.shtml.
            12. *The VLA Expansion Project: Construction Project Book. The Expanded VLA Project Books contains the technical details of the VLA Expansion construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml.
            13. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains a useful source of information on much of the VLA. There are a few hard copies at the VLA and in the DSOC.  Much of this document is now available for downloading.  Note: it does not include any information about the hardware and software specific to the expanded Karl G. Jansky VLA.

            Key Personnel

            Please direct queries to the NRAO Helpdesk; you can expect a response within one business day..  

            If you wish to contact an NRAO staff member directly, you may e-mail them by addressing your message to "first initial last name"@nrao.edu. Thus, you may contact Joanne Astronomer at: "jastrono@nrao.edu". The name is truncated to eight characters.   For questions about telescope time allocation, please e-mail "schedsoc@nrao.edu".

            The listed four-digit numbers are sufficient for calls made from within the SOC.   If you are calling from outside the SOC, dial (+1) 575-835 before dialing the four-digit number. Please also refer to our searchable NRAO Directory

            Editor's Notes

            This Observational Status Summary for the Karl G. Jansky (expanded) VLA is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts.  For questions on the content, or suggestions that would enhance the clarity of this guide, we recommend contacting the NRAO helpdesk.

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            Performance of the VLA during the Next Semester

            Introduction

            This section contains details of the VLA's resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of radio interferometry and synthesis imaging are found elsewhere. In particular, see References 1 and 2, listed in the Documentation.

            Resolution

            VLA Resolution

            The VLA's resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is "blind" to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the VLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/Bmax) are not seen – the resulting image will be smoothed to the resolution of the array. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can partly or fully recover this missing information, which can only be obtained by observing in a more compact array configuration, by using the mosaicing method, or by utilizing data from an instrument such as a large single antenna or a compact array comprising smaller antennas which provides this information.

            Table 2 summarizes the relevant information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power) for the central frequency for each band, and the scale at which severe attenuation of large scale structure occurs.

            These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith.
            Footnotes:
            1. Bmax is the maximum antenna separation, Bmin is the minimum antenna separation, θHPBW is the synthesized beam width (FWHM), and θLAS is the largest scale structure "visible" to the array.
            2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (DnC, CnB, BnA) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., DnB, or CnA) is provided.
            3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θHPBW. The values for snapshots are about 1.3 times the listed values.
            4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two.
            5. For the C configuration an antenna from the middle of the north arm is moved to the central pad "N1". This results in improved imaging for extended objects, but may slightly degrade snapshot performance. Note that although the minimum spacing is the same as in D configuration, the surface brightness sensitivity and image fidelity to extended structure is considerably inferior to that of the D configuration.

            The following figure is a graphical representation of the synthesized beamwidths for natural and robust weighting for the four main array configurations between 1 and 50 GHz. Also available are synthesized beamwidth figures for the low frequency (1–12 GHz) and the high frequency (12–50 GHz) receiver bands.

            A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; the goal of implementing a new digital Pie Town link, now that EVLA construction is complete, remains unfunded.

            Sensitivity

            The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by:

             

             

            where:

            - SEFD is the "system equivalent flux density" (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEFD indicate more sensitive performance. For the VLA's 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62TsysA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band.
            - ηc is the correlator efficiency (~0.93 with the use of the 8-bit samplers).
            - npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in 'RCP' or 'LCP'.
            - N is the number of antennas.
            - tint is the total on-source integration time in seconds.
            - Δν is the bandwidth in Hz.

            Figure 1 shows the SEFDs as a function of frequency used in the VLA exposure calculator for those Cassegrain bands currently installed on VLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 3 gives the SEFD at some fiducial VLA frequencies.

            Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical:

            • For the more commonly-used "robust" weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values.
            • Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 3 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 3; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This primarily affects low frequency observations.
            • Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission, primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band.
            • Losses from the 3-bit samplers.  The VLA's 3-bit samplers incur an additional 10 to 15% loss in sensitivity above that expected -- i.e., the efficiency factor ηc = 0.78 to 0.83.

             

              In general, the zenith atmospheric opacity to microwave radiation is very low - typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable.
              Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz.
              Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. However, the antenna tip capability is currently unavailable for the upgraded VLA -- it is hoped that this will again be available in the next year.
              When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item.
              • Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target.

                Both CASA and AIPS allow the application of elevation-dependent gains and an estimated opacity generated from ground-based weather (e.g., through the tasks gencal and plotweather in CASA, and INDXR in AIPS).

              • Pointing. The SEFD quoted above assumes good pointing. Under calm nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime can be much worse -- occasionally exceeding 1 arcminte, due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 15 mph (7 m/s). Wind speeds near the stow limit (45 mph) will have a similar negative effect at 8 and 15 GHz.
              To achieve better pointing, "referenced pointing" is recommended, where a nearby calibrator is observed in interferometer pointing mode every hour or so. The local pointing corrections thus measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2 - 3 arcseconds (but more typically 5 to 7 arcseconds) if the reference source is within about 15 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees. At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup.
              Use of referenced pointing is highly recommended for all Ku, K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive, and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.3 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution.

              To aid VLA proposers there is an on-line guide to the exposure calculator; the exposure calculator provides a graphical user interface to these equations.

              Special caveats apply for P-band (230 – 470 MHz) observing.  The listed SEFD in table 3 is from an observation taken far from the Galactic plane, where the sky brightness is about 30K.  At this band, Galactic synchrotron emission is very bright in directions near the Galactic plane.  The system temperature increase due to Galactic emission will degrade sensitivity by factors of two to three for observations in the plane, and by a factor of 5 or more at or near the Galactic center.   In addition, the antenna efficiency (currently about 0.31 for 300 MHz) will decline with both increasing and decreasing frequencies from the center of P-band.

              The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by:

              where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼ 1.5λ/Bmax), and S in mJy per beam, the parameter F depends on the synthesized beam, therefore on the array configuration, and has the approximate value F = 190, 18, 1.7, 0.16 for A, B, C, and D configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ΔIm, for S. Note that Equation 2 is a beam-averaged surface brightness; if a source size can be measured the source size and integrated flux density should be used in Equation 2, and the appropriate value of F calculated. In general the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 7 of Reference 1, listed in Documentation.

              For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791):

              where D is the distance to the galaxy in Mpc, and SΔV is the HI line area in units of Jy km/s.

              VLA Frequency Bands and Tunability

              Bands

              For observations taken with the 8-bit samplers,  each receiver can tune to two different frequencies, each 1.024 GHz wide,  within the same frequency band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies, except for the low-band receiver (50 – 500 MHz), which provides linear polarization (X and Y).  Each of these four data streams follows the VLA nomenclature, and are known as IF (for "Intermediate Frequency" channel) "A", "B", "C", and "D". IFs A and B provide RCP (or Y when applicable), IFs C and D provide LCP (or X when applicable).  IFs A and C are always at the same frequency, as are IFs B and D (but note that the A and C IFs  frequency is usually different from the B and D frequency). We normally refer to these two independent data streams as "IF pairs" – i.e., the 'A/C' pair and the 'B/D' pair.   In 8-bit mode, a maximum of 1.024 GHz can be correlated for each IF pair (see the WIDAR Section), for a total maximum bandwidth of  2.048 GHz.  To distinguish this 8-bit system from the 3-bit system, these IF pairs are denoted A0/C0 and B0/D0.

              With the 3-bit samplers, more options are available.  This system provides four (R,L) polarization pairs, each 2048 MHz wide.  The A/C IF pair provides two sampled pairs, labelled A1/C1 and A2/C2, and the B/D IF pair provides two sampled pairs, labelled B1/D1 and B2/D2.

              For more details on the 8-bit and 3-bit samplers see the VLA Samplers section.

              The tuning ranges, along with default frequencies for continuum applications, are given in Table 4 below.

              Table 4: Default frequencies for "continuum" applications
              BandRange18-bit continuum applications (GHz)3-bit continuum applications (GHz)

              (GHz)IF pair A0/C0IF pair B0/D0IF pair A1/C1IF pair A2/C2IF pair B1/D1IF pair B2/D2
              4 m (4) 0.058 – 0.0842 .054 – .086
              90 cm (P) 0.23 – 0.473 0.224 – 0.4803
              20 cm (L) 1.0 – 2.04 1.0 – 1.54 1.5 – 2.04
              13 cm (S) 2.0 – 4.0 2.0 – 3.0 3.0 – 4.0
              6 cm (C) 4.0 – 8.0 4.5 – 5.5 5.5 – 6.5 4.0 – 6.0 6.0 – 8.0
              3 cm (X) 8.0 – 12.0 8.0 – 9.0 9.0 – 10.0 8.0 – 10.0 10.0 – 12.0
              2 cm (Ku) 12.0 – 18.0 13.0 – 14.0 14.0 – 15.0 12.0 – 14.0 14.0 – 16.0 16.0 – 18.0
              1.3 cm (K) 18.0 – 26.5 20.2 – 21.2 21.2 – 22.2 22.0 – 24.0 24.0 – 26.0 18.0 – 20.0 20.0 – 22.0
              1 cm (Ka) 26.5 – 40.0 32.0 – 33.0 31.0 – 32.0 33.0 – 35.0 35.0 – 37.0 29.0 – 31.0 31.0 – 33.0
              0.7 cm (Q) 40.0 – 50.0 40.0 -- 41.0 41.0 – 42.0 44.0 – 46.0 46.0 – 48.0 40.0 – 42.0 42.0 – 44.0
              Notes:
              1.  Listed here are the nominal band edges.  For all bands, the receivers can be tuned to frequencies outside this range, but at the cost of diminished performance.  Contact VLA staff for further information.
              2. The 4-band system is currently under development. Observing time may be requested through the RSRO program.
              3. The default setup for P-band will provide  16 subbands from the A0/C0 IF pair, each 16 MHz wide, to cover the frequency range 224 – 480 MHz. The channel resolution is 125 kHz. 
              4. The default frequency set-up for L-band comprises two 512 MHz IF pairs (each comprising 8 contiguous subbands of 64 MHz) to cover the entire 1 – 2 GHz of the L-band receiver.

              Tuning Restrictions

              In general, for all frequency bands except Ka, if the total span of the two independent IF pairs of the 8-bit system (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For K, Ka and Q bands (the only bands where a span greater than 8 GHz is possible), there are special rules:

              • At Ka band, the low frequency edge of the A0/C0 IF pair must be greater than 32.0 GHz. There is no restriction on the B0/D0 frequency, unless the B0/D0 band overlaps the A0/C0 band when the latter is tuned at or near the 32.0 GHz limit.  In this case, the Observation Preparation Tool (OPT) may not allow the requested frequency setups.  Users wanting to use such a frequency setup are encouraged to contact the NRAO Helpdesk for possible tuning options.
              • At K and Q bands, if the frequency span is greater than 8.0 GHz, the B0/D0 frequency must be lower than the A0/C0 frequency.

              For the 3-bit system the maximum frequency span permitted for the A1/C1 and A2/C2 IF pairs is about 5000 MHz.  The same restriction applies to B1/D1 and B2/D2.  The tuning restrictions given above for the separation and location of the 8-bit pairs A0/C0 and B0/D0 also apply to the 3-bit pairs, with A0/C0 replaced by A1/C1 and A2/C2, and B0/D0 replaced by B1/D1 and B2/D2.

              VLA Samplers

              The VLA is equipped with two different types of samplers, 8-bit, 1GHz bandwidth, and 3-bit, 2GHz bandwidth. The choice depends on your science goals, and on technicalities described below.

              The 8-bit Set consists of four 8-bit samplers running at 2.048 GSamp/sec.  The four samplers are arranged in two pairs, each pair providing 1024 MHz bandwidth in both polarizations.  The two pairs are denoted A0/C0 and B0/D0.  Taken together, the four samplers offer a maximum of 2048 MHz coverage with full polarization. The frequency spans sampled by the two pairs need not be adjacent.  Some restrictions apply, depending on band, as described in the section on Frequency Bands and Tunability.

              The 3-bit Set consists of eight 3-bit samplers running at 4.096 GSamp/sec.  The eight samplers are arranged as four pairs, each pair providing 2048 MHz bandwidth in both polarizations.  Two of these pairs, denoted A1/C1 and A2/C2 cannot span more than 5000 MHz (lower edge of one to the higher edge of the other). The same limitation applies to the second pair, denoted B1/D1 and B2/D2.  The tuning restrictions are described in the section on  Frequency Bands and Tunability.

               

              Which set to use?

              • S, L, and 4/P-band observations, whether line or continuum, should use the 8-bit sampler set.
              • C and X-band continuum observations should use 3-bit samplers in order to exploit the full 4 GHz bandwidth: in spite of the 15% reduction in sensitivity that comes with 3-bit (at equal bandwidth to the 8-bit samplers; see below for details) and the reduced effective bandwidth after removing RFI this still provides superior overall sensitivity. For more details we refer to EVLA memo 166.
              • Ku, K, Ka, and Q band continuum observations should use the 3-bit samplers for maximum bandwidth.
              • Wide-band spectral line searches requiring more than 2 GHz span should use the 3-bit samplers.
              • Spectral-line observations which fit within two (possibly disjoint) 1 GHz bands should use the 8-bit set.
              • Simultaneous continuum and high resolution spectral line observation can use mixed 3-bit and 8-bit samplers. The 3 bit samplers in this case will be set up to deliver the continuum data, while the 8-bit samplers will be for the spectral line data. The mix mode can be used in C and the higher frequency bands.


              Major Characteristics of each Set

              The 8-bit samplers are warranted for observations at 4/P-band, L-band, and S-band. The full analog bandwidth from the receivers fits within the 2048 MHz span covered by the samplers.

              For the 3-bit samplers users need to be aware of the following issues:

              • Sensitivity: compared to the 8-bit system, the sensitivity of  the 3-bit samplers is worse by ~15% (at equal bandwidth). Alternatively, a given continuum noise level, requiring on-source  integration time T with the 8-bit (two bands of 1GHz), requires  0.33T with the 3-bit (4 bands of 2GHz, assuming the bandwidth is available from the front end.)
              • Resonances: each of the eight 3-bit samplers on an antenna has a resonance about 3 MHz wide.  Each resonance is independent of all others, so there is no correlated signal between antennas.  The resonance degrades the spectrum in its narrow frequency range, but has little effect on continuum observing.  Bandpass solutions will be affected, but can be interpolated over. Spectral-line calibration and  images at  the affected frequencies will show significant loss in sensitivity. The resonances are easily seen in autocorrelation spectra, and it is recommended that users, especially spectral-line users, utilize these to locate the compromised frequencies.
              • Amplitude Calibration: The traditional method for both 8- and 3-bit systems is to observe a flux-density calibrator, use self-cal to determine the antenna amplitude calibration factors (gains), and transfer the gains to the phase calibrator and target. For 3-bit samplers this procedure gives results good to  5%, between elevations of 20-70degrees. (Expect worse at the upper edge of Q-band and/or during bad weather). The switched power data can be used to correct for system gain variations and works well for the 8-bit samplers.  For 3-bit samplers, the Pdif depends on the Psum,  i.e. Pdif is non-linear, and its application will bias the resulting visibilities by 5 to 10%.  The origin of this effect is understood, but we have not yet determined how best to compensate for it.  Because of this, we do not recommend use of the 'Psum' and 'Pdif' data to calibrate visibilities from the 3-bit samplers. We do, however, recommend that the 'Requantizer Gains' in the switched power data be applied to remove gain changes. For more information about the switched power, and 'Psum' and 'Pdif, see EVLA memo 145.

               

              Setting up the 8-bit or 3-bit Samplers

              Either set requires an initial scan for each individual LO (frequency) tuning, during which power levels are optimized.

              For the 8-bit system, a 'dummy scan' of 1 minute duration is sufficient for each tuning.  This  is usually done while the antennas are slewing at the start of an observing file, as the pointing direction of the antennas is not critical.

              For the 3-bit system, the requirements are more demanding, see the section on 3-bit set-up. The minimum setup time is 1 minute for each tuning, to adjust the power levels and bandpass slopes across the 2GHz samplers.  These values are retained and applied if the tuning is re-encountered in the same observation.  In addition, every time the LO setup is changed, whether or not it is new (e.g. changing from 8-bit X-band reference pointing back to target),  a scan of 30 seconds is needed to re-set the subband gains (requantizers) in the correlator.  For better amplitude calibration at high frequencies, the 3-bit initial setup should be near the elevation of the target, so do it after the first 8-bit setup described above. For 3-bit observing without 8-bit (e.g. C or X-band, without reference pointing), the power variation with elevation is small so the 3-bit setup can be done at any elevation.

              For settings that use a mix of 3-bit and 8-bit samplers, the guidelines to set up the 3-bit samplers should be followed.


              Other issues

              The overhead for setup of 3-bit samplers can eat into observing time, especially for projects with many different LO settings, and/or sources all over the sky, accompanied by a band change, reference pointing, and requantizer reset for each direction. The impact is most severe for short scheduling blocks.

              Polarization testing conducted so far indicates no degradation of performance by using the 3-bit samplers.

               

              Acknowledgements

              NRAO is grateful to Professor Rob Ivison for supporting the upgrade of some of the 3-bit samplers on the VLA via a grant from the European Research Council.  For observations using the 3-bit samplers between May 2015 and March 2018 we encourage users to include the following text in the Acknowledgments section of their publications:

              "We acknowledge funding towards the 3-bit samplers used in this work from ERC Advanced Grant 321302, COSMICISM."

              Field of View

              Primary Beam

              The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θPB = 45/νGHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS (task PBCOR) and CASA (clean task and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as "mosaicing," in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 in Documentation for details.

              Guidelines for mosaicing with the VLA are given in the Guide to Observing with the VLA

              Chromatic Aberration (Bandwidth Smearing)

              The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array.

              These effects can be parameterized by the product of the fractional bandwidth (Δν/ν0) with the source offset in synthesized beamwidths (θ0HPBW). Table 5 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 5 should be used to determine how much spectral averaging can be tolerated when imaging a particular field.

              Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). Δν/ν0 is the fractional bandwidth; θ0HPBW is the source offset from the phase tracking center in units of the synthesized beam.

              Time-Averaging Loss

              The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parametrized, except for observations at δ = 90°. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with Δν/ν0 replaced by ωeΔtint, where ωe is the Earth's angular rotation rate, and Δtint is the averaging interval.

              For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 (in Documentation) considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 6, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θΔtint)2, where θ is the distance from the phase center and Δtint is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see the section on Time Resolution and Data Rates) in the A and B configurations.

              Note: The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines.

              Non-Coplanar Baselines

              The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations.

              Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the "W-projection" method available in CASA; see EVLA Memo 67 for more details.

              Time Resolution and Data Rates

              The default integration times for the various array configurations are as follows:

              ConfigurationObserving
              Bands
              Default
              integration time
              D, C L S C 5 seconds
              D, C X Ku K Ka Q 3 seconds
              B all 3 seconds
              A all 2 seconds

              Observations with the 3-bit (wideband) samplers must use these integration times. Observations with the 8-bit samplers may use shorter integration times, but these must be requested and justified explicitly in the proposal, and obey the following restrictions:

              Minimum Integration Times and Maximum Data Rates
              Proposal type

              Minimum integration time

              Maximum data rate
              General Observing (GO) and
              Shared Risk Observing (SRO)
              50 msec 25 MB/s (90 GB/hr), or up to
              60 MB/s (216 GB/hr) with additional justification
              Resident Shared Risk Observing (RSRO) < 50 msec > 60 MB/s (216 GB/hr)

              Note that integration times as short as 5 msec and data rates as high as 300 MB/s can be supported for some observing, though any such observing is considered Resident Shared Risk Observing (RSRO). For these short integration times and high data rates there will be limits on bandwidth and/or number of antennas involved in the observation. Those desiring to utilize such short integration times and high data rates should consult with NRAO staff.

              The maximum recommended integration time for any EVLA observing is 60 seconds. For high frequency observations with short scans (e.g., fast switching, as described in Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)), shorter integration times may be preferable.

              Observers should bear in mind the data rate of the VLA when planning their observations. For Nant antennas and integration time Δt, the data rate is:

              Data rate ~ 45 MB/sec × (Nchpol/16384) x Nant × (Nant - 1)/(27×26) / (Δt/1 sec)
              ~ 160 GB/hr × (Nchpol/16384) x Nant × (Nant - 1)/(27×26) / (Δt/1 sec)
              ~ 3.7 TB/day × (Nchpol/16384) x Nant × (Nant - 1)/(27×26) / (Δt/1 sec)
              ...equation (4)

              Here Nchpol is the sum over all subbands of spectral channels times polarization products:

              Nchpol = Σsb Nchan,i x Npolprod,i

              where Nchan,i is the number of spectral channels in subband i, and Npolprod,i is the number of polarization products for subband i (1 for single polarization [RR or LL], 2 for dual polarization [RR & LL], 4 for full polarization products [RR, RL, LR, LL]). This formula, combined with the maximum data rates given above, imply that observations using the maximum number of channels currently available (16384) will be limited to minimum integration times of ~2 seconds for standard observations, and 0.8 seconds for shared risk observations.

              These data rates are challenging for data transfer, as well as data analysis. Currently data may either be downloaded via ftp within the Science Operations Centers, or mailed on hard drives for those not in the same building as the archive. The Archive Access Tool allows some level of frequency averaging to decrease data set sizes before ftp, for users whose science permits; note that the full spectral resolution will be retained in the NRAO archive for all observations.

              Higher time resolutions and data rates are possible in principle but will be considered only through the Resident Shared Risk program.

              Note: The data rate formula given above does not account for the auto-correlations delivered by WIDAR. Precise data rate values can be obtained through the use of the General Observing Setup Tool (GOST), or the Resource Catalog of the Observation Preparation Tool.

              Radio-Frequency Interference

              The very wide bandwidths of the upgraded Very Large Array mean that RFI (radio-frequency interference) will be present in a far larger fraction of VLA observations than in observations made with the old systems.  Considerable effort has gone into making the VLA's new electronics as linear as possible, so that the effects of any RFI will remain limited to the actual frequencies at which the RFI exists.   Non-linear effects, such as receiver saturation, should occur only for those very unlikely, and usually very brief, times when the emitter is within the antenna primary beam.

              RFI is primarily a problem within the low frequency bands (C, S, L, and the low-band system), and is most serious to the D configuration.  With increasing frequency and increasing resolution comes an increasing fringe rate, which is often very effective in reducing interference to tolerable levels.

              The bands within the tuning range of the VLA which are allocated exclusively to radio astronomy are 1400-1427 MHz, 1660-1670 MHz, 2690-2700 MHz, 4990-5000 MHz, 10.68-10.7 GHz, 15.35-15.4 GHz, 22.21-22.5 GHz, 23.6-24.0 GHz, 31.3-31.8 GHz, and 42.5-43.5 GHz. No external interference should occur within these bands.

              RFI seen in VLA data can be internal or external.  Great effort has been expended to eliminate all internally-generated RFI.  Nevertheless, some internal RFI remains, which we are working hard to eliminate.   Nearly all such internally-generated signals are at multiples of 128 MHz.  So far as we know, all such internal signals are unresolved in frequency, and hence will affect only a single channel.

              Radio frequency interference of external origin will be an increasing problem to astronomical observations. Table 7 lists some of the sources of external RFI at the VLA site that might be observed within the VLA's expanded tuning range within L and S bands. Figure 3 shows a raw power cross-power spectrum at L-band. Figure 4 shows a similar plot for the lower half of S-band.

               

               

               

              The three last entries in the table deserve extra discussion.  These are all satellite transmissions, whose severity is a strong function of the angular offset between the particular satellite and the antenna.  It appears that significant degradation can occur if the antennas are within ~10 degrees of the satellite.  The great majority of the satellites are along the 'Clarke Belt' -- the zone of geosynchronous satellites.  As seen from the VLA, this belt is at a declination of about -5.5 degrees.  There are dozens -- probably hundreds -- of satellites 'parked' along this belt, transmitting in many bands:  S, C, Ku, K, and Ka at a minimum.  Observations of sources in the declination rate of +5 to -15 degrees can expect to be significantly degraded due to satellite transmission.  The Sirius digital radio system (and probably the satellites in the 2178 -- 2195 MHz band) comprises three satellites in a 24-hour, high eccentricity orbit with the apogee above the central U.S.  For the Sirius system, the orbit is arranged such that each of the three satellites spends about eight hours near an azimuth of 25 degrees and an elevation of 65 degrees.  The corresponding region in astronomical coordinates is between declinations 50 to 65 degrees, and hour angles between -1 and -2 hours.  Observations at S-band within that area may -- or may not -- be seriously affected.  The most reliable way to judge the seriousness of the satellite emissions is to inspect the switched power table data, particularly in those subbands where there is little RFI.

              VLA staff periodically observes the entire radio spectrum, with the VLA,  from 1.0 through 50.0 GHz with 125 kHz channel resolution to monitor the ever-changing RFI spectrum.  Plots from this program, accompanied with tables of identified sources are available at the RFI section of the VLA Observing Guide. Users concerned about the precise frequencies of strong RFI, and the likelihood of being affected, are encouraged to peruse these plots.

              Although most of the stronger sources of RFI are always present, it is very difficult to reliably predict their effect on observations.  Besides the already noted dependence on frequency and array configuration, there is another significant dependency on sky location for those satellites in geostationary orbit.  For these transmitters, (for example, the frequency range from 3.8 to 4.2 GHz), the effect on observing varies dramatically on the declination of the target source.  Sources near zero declination will be very strongly affected, while observations north of the zenith may well be nearly unaffected, especially at the highest resolutions.

              Also available are total-power plots of all RFI observations made by the interference protection group, from 1993 onwards at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI environment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu.

              The VLA electronics (including the WIDAR correlator) have been designed to minimize gain compression due to very strong RFI signals, so that in general it is possible to observe in spectral windows containing RFI, provided the spectra are well sampled to constrain Gibbs ringing, and spectral smoothing (such as Hanning) is applied.  Both AIPS and CASA provide useful tasks which automatically detect and flag spectral channels/times which contain strong RFI.

              Extracting astronomy data from frequency channels in which the RFI is present is much more difficult. Testing of algorithms which can distinguish and subtract RFI signals from interferometer data is ongoing.

              The 3-bit samplers will be more susceptible to RFI signals than the 8-bit samplers, since the latter have more 'levels' within which these strong signals can be accommodated.  However, since the RFI power at the bands where the 3-bit samplers will most commonly be used (C, X, Ku, K, Ka, and Q) is nearly always less than the total noise power, we do not expect problems when wide-band 3-bit observing is done in these bands.

              Calibration of VLA data when strong RFI is present within a subband can be difficult.  Careful editing of the data, using newly available programs within CASA and AIPS, will be necessary before sensible calibration can be done.  The use of spectral smoothing (typically, Hanning), prior to editing and calibration, is strongly recommended when RFI is present within a subband.

              Identification and removal of RFI is always more effective when the spectral and temporal resolutions are high.  However, the cost of higher spectral and temporal resolution is in database size and, especially, in computing time.  A good strategy is to observe with high resolution, then average down in time and frequency once the editing is completed.

              Subarrays

              The continuum subarray option offers two 1 GHz baseband pairs with the 8-bit samplers in up to 3 subarrays, with the same spectral channel and polarization product options as are available for wideband observing. The setup for each subarray is completely independent, in terms of observing frequency, polarization products, and integration times.

              When using three subarrays, there are some restrictions on the number of antennas in each subarray. The Baseline Board in the correlator treats each set of 4 antennas independently, using a separate column of correlator chips. With 8 such columns, the correlator can handle up to 8x4= 32 antennas. The correlator configuration software requires that a given column not be split across subarrays. This does not matter when using only two subarrays, but forces some subtle restrictions when using three. For instance, one cannot observe with 9 antennas in each of 3 subarrays, because 9 antennas requires three columns (two with 4 antennas each, and one with 1 antenna); three subarrays of 9 antennas each would require 3x3= 9 columns, one more than are actually available. Splitting the array into 10, 9, and 8 antennas is allowed, since the first two subarrays use 3 columns each, while the third uses only two.

              The following table gives four examples of how correlator resources can be split into multiple subarrays. Antennas in each subarray are color-coded: red for subarray 1, green for subarray 2 (if present), and blue for subarray 3 (if present).  The last column gives the number of antennas in each subarray (e.g., in the setup shown in the first row, subarray 1 has 10 antennas; subarray 2 has 9 antennas; and subarray 3 has 8 antennas). In all cases a total of 27 antennas are used.  [The columns are numbered in reverse order (C7 to C0) to match the numbering scheme used on the actual Baseline Boards.]

              Some Possible Subarray Options
              Number of antennas correlated using each Baseline Board columnNumber
              C7 C6 C5 C4 C3 C2 C1 C0 of antennas
              4 4 2 4 4 1 4 4 10 + 9 + 8
              4 4 4 2 4 4 4 1 14 + 13
              4 4 3 4 4 3 3 2 11 + 11 + 5
              4 4 4 4 4 4 3 -- 27

              Positional Accuracy & Astrometry

              Summary: The position of a target can be determined to a small fraction of the synthesized beam, limited by atmospheric phase stability, the proximity of an astrometric calibrator, the calibrator-source cycle time, and the SNR on target.

              In preparation for observing, the a-priori position must be known to within the antenna primary beam, except perhaps for mosaicing observations. In the special case of using the phased VLA as a VLBI element, the a-priori position must be accurate to within the synthesized beam of the array.

              In post-processing, target positions are typically determined from an image made after phase calibration, i.e. correcting the antenna and atmospheric phases as determined on the reference source. The accuracy of the calibration determines the  accuracy of the positions in the image. (Note that phase self-calibration imposes the assumed position of the model, i.e.,  makes  the position indeterminate. Therefore, an absolute position cannot be determined after self-calibration, but relative positions between features within a self-calibrated image are valid.)

              It may help to think of astrometry in 2 steps, narrow and wide-field.

              In narrow-field astrometry, the target is close to the phase tracking center and the antennas nod every few minutes between the target and a calibrator. Under good conditions of phase stability, accurate antenna positions, (so-called 'baselines'), a strong target, a close calibrator with accurately known position, and rapid switching, the accuracy can approach 1-2% of the synthesized beam, with a floor of ~2 mas. Even under more typical conditions, 10% of the beam is readily achieved.

              Astrometric calibrators are marked 'J2000  A' in the VLA calibrator list, and have an accuracy of ~2 mas. Other catalogs from the USNO and the VLBA are also useful, but offsets may exist between the VLA and VLBA centroids, arising from extended structure in the particular source, and the different resolutions of the arrays.

              For studies of proper motion and parallax, the absolute accuracy of a calibrator may be less important than its stability over time. Close or in-beam calibrators with poor a-priori positions can be used, and tied to the ICRF reference frame in the same or separate observations.

              Phase stability can be assessed in real time from the Atmospheric Phase Interferometer (API) at the VLA site, which uses observations of a geostationary satellite at ~12GHz. Dynamic scheduling uses the API data to run a project under suitable conditions, specified by the user. Note that  VLBI projects using the phased VLA will typically be fixed date, not dynamically scheduled.

              The widefield case is to determine the positions of targets within the primary beam, referenced to a calibrator within the beam or close by. In addition to the previous effects, there are distortions as a function of position in the field, from small errors in the Earth orientation parameters (EOP) used at correlation time, differential aberration, and phase gradients across the primary beam.  With no special effort, the errors build up to roughly ~1 synthesized beam at a separation of ~10^4 beams from the phase tracking center.  Not all these errors are fully understood, and accurate recovery of positions over the full primary beam in the wideband, widefield  case is a research area. These effects are handled somewhat differently in the post-processing packages. Check with VLA staff for more details.

              Limitations on Imaging Performance

              Image Fidelity

              Image fidelity is a measure of the accuracy of the reconstructed sky brightness distribution. A related metric, dynamic range, is a measure of the degree to which imaging artifacts around strong sources are suppressed, which in turn implies a higher fidelity of the on-source reconstruction.

              With conventional external calibration methods, even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred.  The limiting factor is most often the effective phase stability of the telescope due to atmospheric/ionospheric fluctuations, although pointing errors and changes in atmospheric opacity can also be a limiting factor.  If a good model of the sky brightness distribution exist (e.g. use of compact structures of sufficient strength, though a good model of resolved sources in the field of view may also be used), standard self-calibration can be counted on to improve the images.  At low frequencies where the dominant phase error is due to ionospheric plasma density fluctuations, more advanced techniques may be required to account for change of ionospheric phase across the field of view.  Dynamic ranges in the thousands to hundreds of thousands can be achieved using these techniques, depending on the underlying nature of the errors. With the new WIDAR correlator and its much greater bandwidths and much higher sensitivities, self-calibration methods can be extended to observations of sources with much lower flux densities than very possible with the old VLA.

              The choice of image reconstruction algorithm also affects the correctness of the on-source brightness distribution. The CLEAN algorithm is most appropriate for predominantly point-source dominated fields. Extended structure is better reconstructed with multi-resolution and multi-scale algorithms. For high dynamic ranges with wide bandwidths, algorithms that model the sky spectrum as well as the average intensity can yield more accurate reconstructions.

              Invisible Structures

              An interferometric array acts as a spatial filter, so that for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Image reconstruction algorithms such as multi-resolution and multi-scale CLEAN can help to reduce or eliminate these negative bowls, but care must be taken in choosing appropriate scale sizes to work with.

              Table 2 gives the largest scale visible to each configuration/band combination.

              Poorly Sampled Fourier Plane

              Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in Documentation.

              Sidelobes from non-Deconvolved Sources

              At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task.  Sidelobe confusion is a strong function of observing band -- affecting most strongly L and P-band observations.  It is rarely a significant problem for observations at frequencies above 4 GHz.

              Sidelobes from Strong Sources

              An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere.

              Wide Field Imaging

              Wide-field observing refers primarily to the non-coplanar nature of the VLA when observing in non-snapshot mode. At high angular resolutions and low frequencies, standard imaging methods will produce artifacts around sources away from the phase center.  Faceted imaging (AIPS, CASA) and w-projection (CASA) techniques can be used to solve this problem.

              Another aspect of wide-field observing is the accurate representation of primary beam patterns, and their use during imaging. This is relevant only for very high dynamic ranges ( > 10000 ) or when there are very strong confusing sources at and beyond the half-power point of the primary beam.  This problem is worse with a wide-band instrument simply because the size of the primary beam (and the radius at which the half-power point occurs) varies with frequency, while there is also increased sensitivity out to a wider field of view. Work is under way to commission algorithms that deal with these effects by modeling and correcting for frequency-dependent and rotating primary beams per antenna, during imaging. Please note, however, that most advanced methods will lead to a significant increase in processing time, and may not always be required. Therefore, in the interest of practicality, they should be used only if there is evidence of artifacts without these methods.

              Finally, all of the above effects come into play for mosaicing, another form of wide-field imaging in which data from multiple pointings are combined during or after imaging.

              Wide-band Imaging

              The very wide bandpasses provided by the upgraded Very Large Array enable imaging over 2:1 bandwidth ratios -- at L, S, and C bands, the upper frequency is twice that of the lower frequency.   It is this wide bandwidth which enables sub-microJy sensitivity.

              In many cases, where the observation goal is a simple detection, and there are no strong sources near to the region of interest, standard imaging methods that combine the data from all frequencies into one single image (multi-frequency-synthesis) may suffice.  This is because the wide-band system makes a much better synthesized beam -- especially for longer integrations -- than the old single-frequency beam, thus considerably reducing the region of sky which is affected by incorrect imaging/deconvolution.  A rough rule of thumb is that -- provided a strong source is not adjacent to the target zone -- if the necessary dynamic range in the image is less than 1000:1, (i.e., the strongest source in the beam is less than 1000 times higher than the noise), a simple wide-band map may suffice.  

              For higher dynamic ranges, complications arise from the fact that the brightness in the field of view dramatically changes as a function of frequency, both due to differing structures in the actual sources in the field of view, and due to the attenuation of the sources by the primary beam.  One symptom of such problems is the appearance of radial spokes around bright sources, visible above the noise floor, when imaged as described above.  

              The simplest solution is to simply make a number of maps (say, one for each subband), which can then be suitably combined after correction for the primary beam shape. But with up to 64 subbands available with the VLA's new correlator, this is not always the optimal approach.  Further, images at all bands must be smoothed to the angular resolution at the lowest frequency before any spectral information can be extracted, and with a 2:1 bandwidth the difference in angular resolution across the band will be significant.  

              A better approach is to process all subbands simultaneously, utilizing software which takes into account the possibility of spatially variant spectral index and curvature, and knows the instrumentally-imposed attenuation due to the primary beam. Such wideband imaging algorithms are now available within CASA as part of the clean task, and work is under way to integrate them fully with wide-field imaging techniques.

              Calibration and Flux Density Scale

              The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators, and is also available within the Observation Preparation Tool.

              Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration - consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, S, C, X, Ku, K, Ka, and Q bands.

              Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the VLA cannot accurately measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards, as described in the next paragraph. Table 8 shows the flux densities of these sources in January 2012 at the standard VLA bands.  The flux density scale for the VLA, from 1 through 50 GHz, is based on emission models of the planet Mars, which is then calibrated to the CMB dipole using WMAP (Wilkinson Microwave Anisotropy Probe) observations (see Perley and Butler, 2013, for details).  The source 3C286 (=J1331+3030) is known to be non-variable, and has thus been adopted as the prime flux density calibrator source for the VLA. The adopted polynomial expression for the spectral flux density for 3C286 is:

              \[\log(S) = 1.2515 - 0.4605 \log(f) - 0.1715 \log^2(f) + 0.0336 \log^3(f)\]

              where S is the flux density in Jy, and f is the frequency in GHz.

              The absolute accuracy of our flux density scale is estimated to be about 2%.   With care, the internal accuracy in flux density bootstrapping is better than 1% at all bands except Q-band, where pointing errors limit bootstrap accuracy to perhaps 3%.   Note that such high internal accuracies are only possible in long-duration observations where the antenna gains curves and atmospheric opacity can be directly measured, and where there is good elevation overlap between the target source(s) and the flux density standard calibrator.

              Table 8: Flux densities (Jy) of Standard Calibrators for January 2012
              Source

              1465 MHz

              2565 MHz
              4885 MHz 8435 MHz 14965 MHz 22460 MHz 36435 MHz43340 MHz
              3C48 = J0137+3309 15.56 9.80 5.39 3.14 1.77 1.19 0.73 0.63
              3C138 = J0521+1638 8.71 6.17 4.02 2.78 1.89 1.46 1.03 0.92
              3C147 = J0542+4951 21.85 13.75 7.59 4.49 2.59 1.77 1.10 0.94
              3C286 = J1331+3030 14.90 10.03 7.34 5.09 3.39 2.52 1.75 1.53
              3C295 = J1411+5212 22.15 12.95 6.41 3.34 1.62 0.957 0.507 0.403
              NGC7027 1.62 3.59 5.38 5.79 5.62 5.42 5.18 5.04

              The sources 3C48, 3C147, and 3C138 are all slowly variable.  VLA staff monitor these variations on timescale of a year or two, and suitable polynomial coefficients are determined for them which should allow accurate flux density bootstrapping. These coefficients are updated approximately every other year, and are used in the AIPS task SETJY and in the CASA task setjy.

              The VLA antennas have elevation-dependent gain variations which are important to account for at the four highest-frequency bands.  Gain curves are determined by VLA staff, and the necessary corrections are applied to the visibility data when these data are downloaded from the archive. In addition to this, atmospheric opacity will also cause an elevation-dependent gain which is particularly notable at these four highest frequency bands.   At the current time, we do not have an atmospheric opacity monitoring procedure, so users should utilize the appropriate tasks available in both AIPS and CASA to estimate and correct for the opacity using ground-based weather data.  Correction of these gain dependencies, plus regular calibration using a nearby phase calibrator, should enable good amplitude gain calibration for most users.  Note that extraordinary attenuation by clouds can only be (approximately) corrected for by regular observation of a nearby calibrator.

              A better procedure for removing elevation gain dependencies uses the AIPS task ELINT.  This task will generate a 2nd order polynomial gain correction utilizing your own calibrator observations.  This will remove both the antenna and opacity gain variations, and has the decided advantage of not utilizing opacity models or possibly outdated antenna gain curves.  Use of this procedure is only practical if your observations span a wide range in elevation.

              By far the most important gain variation effect is that due to pointing.  Daytime observations on sunny days can suffer pointing errors of up to one arcminute (primarily in elevation).  This effect can be largely removed by utilizing the 'referenced pointing' procedure.  This determines the pointing offset of a nearby calibrator, which is then applied to subsequent target source observations.  It is recommended that this local offset be determined at least hourly, utilizing an object within 15 degrees of the target source -- preferentially at an earlier HA.  Studies show that the maximum pointing error will be reduced to about 7 arcseconds, or better.  VLA staff continue to work on improving this essential methodology.

              The VLA's post-amplifiers are not temperature stabilized, and exhibit significant gain changes between night and day, particularly at the four highest frequency bands.  Changes as large as 30% have been seen between night and day in calm, clear conditions!  These gain changes (and others caused by possible changes in attenuator settings) are monitored and will be removed with excellent accuracy by application of the internal calibration signal, whose results are recorded in the switched power table (SY table, in AIPS).  These corrections are not applied by default -- users who wish to correct for these gain changes must utilize the appropriate tasks in AIPS or CASA.   For the most accurate flux density bootstrapping, this table must be applied to the visibility data before calibration.  Gain bootstrapping better than 1% can be accomplished for the 8-bit sampler system after application of the switched power data.  For the 3-bit system there is an additional complication, as the values of the switched power data are sensitive to the total power, as well as the system gain.  VLA staff are currently working on a methodology to remove the total power dependency.  Not applying the switched power data will reduce bootstrapping accuracy to perhaps 10%, and possibly worse, if the observation of the flux density calibrator is not close in time to the local phase/amplitude calibrator.

              Complex Gain Calibration

              General Guidelines for Gain Calibration

              Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is difficult to define simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below.

              • Under decent conditions (no thunderstorms or ionospheric storms) tropospheric effects dominate at frequencies higher than about 4 GHz, ionospheric effects dominate at frequencies lower than about 4 GHz.
              • Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the C and D configurations at L and S bands, and in the D configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes - a couple of times per hour is generally sufficient.
              • If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale.  The enhanced sensitivity of the VLA guarantees, for full-band continuum observations, that every field will have enough background sources to enable phase self-calibration at L and S bands.  At higher frequencies, the background sky is not sufficient, and only the flux of the target source itself will be available.
              • The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an "S" code in the calibrator database, and a more distant calibrator with a "P" code, the nearby calibrator is usually the better choice.  A detailed description of calibrator codes is available in the Key to the calibrator list.
              • In clear and calm conditions, most notably in the summer, phase stability often deteriorates dramatically after about 10AM, due to small-scale convective cells set up by solar heating.  Observers should consider a more rapid calibration cycle for observations between this time and a couple hours after sundown.
              • At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See Rapid Phase Calibration and the Atmospheric Phase Interferometer (API).
              • Use the figure below to estimate how much time is minimally needed for each gain calibrator scan.  For instance, a 1 Jy calibrator and 4 MHz total bandwidth requires at least 30 seconds on source

               

              Minimum time required on a gain calibrator scan as a function of bandwidth and calibrator flux for the rather extreme case of upper Q-band.  Durations derived from this plot will definitely be sufficient for all other bands.

              Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)

              For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used for very small source-calibrator separations. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the old VLA this was known as "fast-switching." For the upgraded VLA it is just a loop of source-calibrator scans with short scan length. This technique "stops" tropospheric phase variations at an "effective" baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest VLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Note, however, that for the compact D-configuration, and a typical wind velocity, this "effective" baseline length is the same as, or larger than, the longest baseline in the array, and it is not worth the increased overhead of short cycle times. Under these circumstances it is sufficient to calibrate every 5-10 minutes to track the instrumental changes. The fast switching technique will also not work in bad weather (such as rain showers, or when there are well-developed convection cells - most notably, thunderstorms). It is also important to specify correctly the required tropospheric phase stability as measured by the Atmospheric Phase Interferometer at observe time (see below).

              Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from References 9 and 10 in Documentation.  See also the High Frequency Observing guide for additional recommendations on observing at high frequencies.

              An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data are heavily used for the dynamic scheduling of the VLA.

              Characteristic seasonal averages are represented in Table 9 below:

              Note: day indicates sunrise to sunset values; night indicates sunset to sunrise values.

              Polarization

              For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The phase calibrator chosen for the observations can also double as a polarization calibrator provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15deg to 50deg for a few hour track). The minimum condition that will enable accurate polarization calibration from a polarized source (in particular with unknown polarization) is three observations of a bright source spanning at least 60 degrees in parallactic angle (if possible schedule four scans in case one is lost). If a bright unpolarized unresolved source is available (and known to have very low polarization) then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for objects small compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission. 3C48 also can be used for this at frequencies of ~3 GHz and higher, or 3C147 at frequencies abover ~10 GHz.  The table below shows the measured fractional polarization and intrinsic angle for the linearly polarized emission for these four sources in December 2010.  Note that 3C138 is variable -- the polarization properties are known to be changing significantly over time, most notably at the higher frequencies.  See Perley and Butler (2013b) for details.

              More information on polarization calibration strategy can be found in the VLA Observing Guide.

               

              Freq.3C48Pol3C48Ang3C138Pol3C138Ang3C147Pol3C147Ang3C286Pol3C286Ang
              GHz % Deg. % Deg. % Deg. % Deg.
              1.05 0.3 25 5.6 -14 <.05 - 8.6 33
              1.45 0.5 140 7.5 -11 <.05 - 9.5 33
              1.64 0.7 -5 8.4 -10 <.04 - 9.9 33
              1.95 0.9 -150 9.0 -10 <.04 - 10.1 33
              2.45 1.4 -120 10.4 -9 <.05 - 10.5 33
              2.95 2.0 -100 10.7 -10 <.05 - 10.8 33
              3.25 2.5 -92 10.0 -10 <.05 - 10.9

              33

              3.75 3.2 -84 - - <.04 - 11.1 33
              4.50 3.8 -75 10.0 -11 0.1 -100 11.3 33
              5.00 4.2 -72 10.4 -11 0.3 0 11.4 33
              6.50 5.2 -68 9.8 -12 0.3 -65 11.6 33
              7.25 5.2 -67 10.0 -12 0.6 -39 11.7 33
              8.10 5.3 -64 10.4 -10 0.7 -24 11.9 34
              8.80 5.4 -62 10.1 -8 0.8 -11 11.9 34
              12.8 6.0 -62 8.4 -7 2.2 43 11.9 34
              13.7 6.1 -62 7.9 -7 2.4 48 11.9 34
              14.6 6.4 -63 7.7 -8 2.7 53 12.1 34
              15.5 6.4 -64 7.4 -9 2.9 59 12.2 34
              18.1 6.9 -66 6.7 -12 3.4 67 12.5 34
              19.0 7.1 -67 6.5 -13 3.5 68 12.5 35
              22.4 7.7 -70 6.7 -16 3.8 75 12.6 35
              23.3 7.8 -70 6.6 -17 3.8 76 12.6 35
              36.5 7.4 -77 6.6 -24 4.4 85 13.1 36
              43.5 7.5 -85 6.5 -27 5.2 86 13.2 36

              High sensitivity linear polarization imaging may be limited by time dependent instrumental polarization, which can add low levels of spurious polarization near features seen in total intensity and can scatter flux throughout the polarization image, potentially limiting the dynamic range. Preliminary investigation of the EVLA's new polarizers indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible.

              The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun. Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing.

              Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive - 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information. Ionosphere correction can also be performed in CASA using the task gencal; consult the calibration chapter of CASA Cookbook for more details.

              VLBI Observations

              The VLA can participate in VLBI observations.  This is only allowed in phased array mode (single dish is only available through the VLBA Resident Shared Risk Observing program) with restricted WIDAR correlator configurations.  Also note that currently P-band cannot be phased.  For more details see the VLBI at the VLA documentation.  In phased array mode the program TelCal derives the antenna-based delay and phase corrections needed for antenna phasing in real time.  This correction is applied to the antenna signals before they are summed, requantized to 2-bits, and recorded in VDIF format on the Mark5C disk at the VLA site.  The disk(s) are then transported to Socorro, NM and correlated on the DiFX correlator with other VLBI stations which participated in the observation.  Standard VLA data, i.e., correlations between VLA antennas,  are also archived in the NRAO science data archive.

              Snapshots

              The two-dimensional geometry of the VLA allows a snapshot mode whereby short observations can be used to image relatively bright unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest.

              Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm and longer wavelengths in the D configuration. For instance, at 20 cm a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. The deconvolution of the data is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing "Faint Images of the Radio Sky at Twenty-centimeters survey (FIRST, http://www.cv.nrao.edu/first/)" (B configuration), or the NRAO VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) (D configuration, all-sky) surveys.

              Shadowing and Cross-Talk

              Observations at low elevation in the C and D configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA tasks importasdm and importevla) when using the default inputs. AIPS task UVFLG can be used to flag VLA data based on shadowing, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays. The CASA task flagdata can also be used to flag data based on shadowing.

              Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. This effect is important at low frequencies in compact configurations.  Careful examination of the visibilities is necessary to identify and remove this form of interference. The affected data would show time-variable high-amplitude points.

              Combining Configurations and Mosaicing

               

              Any single VLA configuration will allow accurate imaging of a range of spatial scales determined by the shortest and longest baselines. For extended and structured objects, it may be required to obtain observations in multiple array configurations. It is advisable that the frequencies used be the same for all configurations to be combined. The ideal combination of arrays results in a uv-plane with all cells equally filled by uv-points. To first order, this can be achieved by using the beam sizes of the individual arrays to inversely scale the on-source integration time. This approach is equivalent to achieving the same surface brightness sensitivity for all arrays on all scales. For the VLA, observations in the different configurations generate beam sizes that decrease by factors of 3, i.e. C configuration generates a 3 times smaller beam than D configuration, B 3 times smaller than C, and A three times smaller than B. Thus, on-source integrations would increase by about an order of magnitude between each array. Such a drastic increase is very expensive and, in fact, not necessary since some spatial scales are common to more than a single array, which is equivalent to some uv-cells being filled more than others. The best way to fill the uv plane depends on many factors, such as declination of the source, LST time of the observation, and bandwidth.

              For the VLA, experience shows that a factor of about 3 in on-source integration time for the different array configurations works well for most experiments.

              E.g., 20min on-source time in D, 1h in C, 3h in B, and 9h in A should produce a decent map. Using large bandwidths and multi-frequency synthesis will broaden all uv tracks radially and one may need even less array configurations or shorter integration times between the different arrays.

              Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing.  The VLA has no limit on the number of pointings for each mosaic. Typically hexagonal, rectangular, or individual pointing patterns are used and the overlap regions will result in an improved rms over each individual pointing. Given the many, potentially short observations, it is important to obey the data rate limits outlined in the Time Resolution and Data Rates Section. On-the-fly (OTF) mosaics, i.e. dumping the data while moving the telescopes across the source, is also available.

              Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 (Documentation) for more information.

              Guidelines for mosaicing with the VLA are given in the VLA Observing Guide.

              Pulsar Observing

              The VLA can be used for two kinds of pulsar observing: phase-binning using the WIDAR correlator, and using the phased-array for pulsar dedispersion (either search or fold mode). Either of these types of observing are considered Resident Shared Risk (RSRO), and close collaboration with NRAO staff is required for their use.

              Complex Gain Calibration

              General Guidelines for Gain Calibration

              Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is difficult to define simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below.

              • Under decent conditions (no thunderstorms or ionospheric storms) tropospheric effects dominate at frequencies higher than about 4 GHz, ionospheric effects dominate at frequencies lower than about 4 GHz.
              • Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the C and D configurations at L and S bands, and in the D configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes - a couple of times per hour is generally sufficient.
              • If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale.  The enhanced sensitivity of the VLA guarantees, for full-band continuum observations, that every field will have enough background sources to enable phase self-calibration at L and S bands.  At higher frequencies, the background sky is not sufficient, and only the flux of the target source itself will be available.
              • The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an "S" code in the calibrator database, and a more distant calibrator with a "P" code, the nearby calibrator is usually the better choice.  A detailed description of calibrator codes is available in the Key to the calibrator list.
              • In clear and calm conditions, most notably in the summer, phase stability often deteriorates dramatically after about 10AM, due to small-scale convective cells set up by solar heating.  Observers should consider a more rapid calibration cycle for observations between this time and a couple hours after sundown.
              • At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See Rapid Phase Calibration and the Atmospheric Phase Interferometer (API).
              • Use the figure below to estimate how much time is minimally needed for each gain calibrator scan.  For instance, a 1 Jy calibrator and 4 MHz total bandwidth requires at least 30 seconds on source

               

              Minimum time required on a gain calibrator scan as a function of bandwidth and calibrator flux for the rather extreme case of upper Q-band.  Durations derived from this plot will definitely be sufficient for all other bands.

              Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)

              For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-Calibrator observing cycles as short as 40 seconds can be used for very small source-calibrator separations. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the old VLA this was known as "fast-switching." For the upgraded VLA it is just a loop of source-calibrator scans with short scan length. This technique "stops" tropospheric phase variations at an "effective" baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest VLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Note, however, that for the compact D-configuration, and a typical wind velocity, this "effective" baseline length is the same as, or larger than, the longest baseline in the array, and it is not worth the increased overhead of short cycle times. Under these circumstances it is sufficient to calibrate every 5-10 minutes to track the instrumental changes. The fast switching technique will also not work in bad weather (such as rain showers, or when there are well-developed convection cells - most notably, thunderstorms). It is also important to specify correctly the required tropospheric phase stability as measured by the Atmospheric Phase Interferometer at observe time (see below).

              Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from References 9 and 10 in Documentation.  See also the High Frequency Observing guide for additional recommendations on observing at high frequencies.

              An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase using an interferometer comprising two 1.5 meter antennas separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data are heavily used for the dynamic scheduling of the VLA.

              Characteristic seasonal averages are represented in Table 9 below:

              Note: day indicates sunrise to sunset values; night indicates sunset to sunrise values.

              The WIDAR Correlator

              Introduction

              The correlator configurations offered for general observing may be divided into three basic modes: wideband, spectral line, and subarrays. Note that the possible setups are also subject to the integration time and data rate restrictions outlined in the section on Time Resolution and Data Rates. The possibilities and restrictions are embodied in the General Observing Setup Tool (GOST) and in the Resources section of the Proposal Submission Tool (PST), which must be used to define the correlator configuration for General Observing (GO) and Shared Risk Observing (SRO) proposals.

              Note that phased array configurations are only allowed as part of VLBI experiments (see the section on VLBI Observations) or as Resident Shared Risk observations.

              Wideband Observing

              The wideband observing setups provide the widest possible bandwidth for a given observing band, with channel spacing depending on the number of polarization products as listed in the following table:

              Wideband & Subarray Correlator Options (all but P- and L-bands)

              Polarization products Channel spacing
              Full (RR, RL, LR, LL) 2   MHz
              Dual (RR, LL) 1   MHz
              Single (RR or LL) 0.5 MHz

              8-bit wideband setups are available for all observing bands, providing a total of 2 GHz of bandwidth per polarization (1 GHz per polarization at L-band, and 256 MHz per polarization at P-band).  3-bit setups are available for all bands above S-band, providing total bandwidths per polarization of 4 GHz (C/X bands), 6 GHz (Ku band), or 8 GHz (K/Ka/Q bands). In all cases but P- and L-band each of the subbands is 128 MHz wide. At L-band the default is 64 MHz/subband, yielding channels twice as narrow as those listed in the table above, while at P-band the default is 16 MHz/subband,  resulting in 125 kHz channel spacing.

              In many frequency bands the total processed bandwidth is less than that delivered by the front-end. In those cases the observer may independently tune two 1 GHz baseband pairs when using the 8-bit samplers, or four 2 GHz baseband pairs when using the 3-bit samplers, or choose to have a mix 8-bit and 3-bit samplers. The tuning restrictions are described in the section on VLA Frequency Bands and Tunability, and the 8-bit and 3-bit samplers are described in the section on VLA Samplers.

              Spectral Line Observing

              Basebands and Subbands

              Currently observers have access to very flexible correlator configurations using up to 64 subbands in up to 4 basebands sampled with the 8-bit and/or the 3-bit samplers.  These capabilities may be summarized as follows:

              • Two 1 GHz baseband pairs using the 8-bit samplers, or four 2 GHz baseband pairs using the 3-bit samplers, independently tunable within the limits outlined in the section on VLA Frequency Bands and Tunability. The 8-bit baseband pairs are referred to as A0/C0 and B0/D0, while the 3-bit samplers are A1/C1, A2/C2, B1/D1, and B2/D2.  The AC/BD nomenclature corresponds to that of the IF pairs in the pre-expansion VLA.
              • Up to 16 subband pairs (spectral windows) in each 3-bit baseband pair, and up to 32 subbands in each 8-bit baseband pair with a total of up to 64 subbands in any correlator configuration
                • Tuning, bandwidth, number of polarization products, and number of channels can be selected independently for each subband
                • All subbands must share the same integration time
                • No part of a subband can cross a 128 MHz boundary
                • Subband bandwidths can be 128, 64, 32, ..., 0.03125 MHz (128 / 2n, n=0, 1, ..., 12)
              • The sum over subbands of channels times polarization products is limited to 16,384 (without recirculation).
                • These may be spread flexibly over subbands and polarization products, in multiples of 64: 64, 128, 192, 256, 384, ..., 16384 cross-correlation products.
                • Recirculation may be used to increase the number of channels per subband for subbands narrower than 128 MHz, Baseline Board stacking may be used to increase the number of channels per subband for setups requiring less than 64 subbands.
                • Assigning many channels to a given subband may reduce the total bandwidth and/or the total number of subbands available.

              The remainder of this section discusses the various limitations in more detail, including some examples to show how they come up in practice.

              Subband tuning restrictions

              Each subband may be placed anywhere within a baseband, with the caveat that no subband may cross a 128 MHz boundary. Mathematically:

              νBB0 + n*128 MHz <= νsbLow <= νsbHigh <= νBB0 + (n+1)*128 MHz

              where:

              νBB0 the lower frequency edge of the baseband;
              n= 0, 1, ..., 7 (, ..., 15) (i.e., any integer between 0 and 7 for 8-bit, between 0 and 15 for 3-bit);
              νsbLow                                   
              the lower edge of the subband (i.e., the subband center frequency minus half the subband bandwidth);
              νsbHigh                                 
              the upper edge of the subband (i.e., the subband center frequency plus half the subband bandwidth).

              So for example, if the baseband were tuned to cover 10000-11024 MHz, one could place a 64 MHz subband to cover 10570-10634 MHz, but not to cover 10600-10664 MHz (because that would cross the 128 MHz boundary at 10640 MHz). Note in particular that the center of a baseband is a boundary and no line should be observed at the baseband center.

              The figure below illustrates these restrictions:

              Correlator configuration figure: bandpass8jul12.png

              The black curve shows the analog filter response for an 8-bit baseband covering 1024 MHz, starting at νBB0. The dashed blue vertical lines show the 128 MHz boundaries; no subband can cross those boundaries.  128 MHz subbands are thus constrained to cover a region between two of those boundaries, and no finer tuning is possible.  Narrower subbands, like the 64 MHz subband shown here in red, can be shifted around arbitrarily within one of the 128 MHz "slots", but cannot cross any of these boundaries. (The dotted vertical red lines show the boundaries of the 64 MHz subband, while the solid curve shows an illustrative line within the subband.)

              The analog filter shape defining the baseband rolls off severely at one edge of the baseband, so the 128 MHz "slot" at that edge has reduced sensitivity.  The baseband edge is at the lowest sky frequency in the baseband when using upper sideband, and at the highest sky frequency in the baseband when using lower sideband.

              Subband bandwidths & the digital filter response

              The bandwidth for each subband may be selected independently, and can be any of 128/2n MHz, for n= 0, 1, ..., 12: 128, 64, 32, 16, 8, 4, 2, or 1 MHz, or 500, 250, 125, 62.5, or 31.25 kHz.

              The usable portion of the subband is set by three effects. First, as discussed above, the analog filters which define the baseband are not perfect, leading to lower sensitivity in the 128 MHz near the baseband edge for the 8-bit samplers

              Second, because the digital filters are not infinitely sharp, the rejected sideband leaks in at both edges of the subband. This leads to additional (aliased) noise, with a factor ~2 increase in the noise at the subband edges, dropping to a few per cent within a few per cent of the subband edge. The precise filter shape and noise increase is a complex but predictable function of the subband bandwidth and the subband tuning.

              The third effect stems from the offset frequencies used for sideband rejection in the WIDAR correlator. The local oscillators at the individual antennas are tuned to slightly different frequencies, with those offsets taken out in the correlator. This means that each antenna observes a slightly different sky frequency, and thus some baselines will not give an interesting correlation near one edge of the subband. The maximum frequency shift is currently set to 32*f0, with the fundamental f0 being set to f0= max(25.6 kHz*sbBW/128 MHz, 100 Hz). Here sbBW is the smallest subband bandwidth within the baseband. For the wider subband bandwidths the maximum frequency shift corresponds to <1% of that bandwidth, but for narrower subbands the effect can be severe. For instance, a 31.25 kHz subband has f0= 100 Hz, and a maximum frequency shift of 3.2 kHz -- 10% of the subband may be lost on some baselines.

              Spectral channels and polarization products

              Each subband (without recirculation enabled) can have a different number of channels and polarization products, subject to two limitations:

              1. For the ithsubband, the number of spectral channels can be:
                • 64 nBlBP,i with full polarization products (RR,RL,LR,LL)
                • 128 nBlBP,i with dual polarization products (RR,LL)
                • 256 nBlBP,i with a single polarization product (RR or LL)
                Here nBlBP,i= 1, 2, 3, 4, 5, ..., 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.
              2. The sum over all subbands of nBlBP,i must be less than or equal to 64, the number of Baseline Board pairs in the correlator. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64x256= 16,384 (without recirculation).

              Baseline Boards are the boards in the WIDAR correlator where the actual cross-multiplications are done. There are 128 Baseline Boards arranged as 64 Baseline Board pairs (BlBPs).  The limitations given here correspond to the capabilities of the individual boards, and the finite number of boards the correlator has.

              Limitation #1 corresponds to the following table of the options for subband bandwidth and spectral resolution when using nBlBP Baseline Board pairs for a subband:

              Subband Bandwidth and Spectral Resolution Options (without recirculation)
              Subband bandwidth &
              total velocity coverage
              Full polarization products
              (RR, RL, LR, LL)
              64nBlBP spectral channels

              Channel spacing:
              Dual polarization products
              (RR, LL)
              128nBlBP spectral channels
              Channel spacing:
              Single polarization product
              (RR or LL)
              256nBlBP spectral channels

              Channel spacing:
              128 MHz 38400/νGHz km/s 2000/nBlBP kHz 600/nBlBPGHz km/s 1000/nBlBP kHz 300/nBlBPGHz km/s 500/nBlBP kHz 150/nBlBPGHz km/s
              64 19200 1000 / nBlBP 300 / nBlBP 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP
              32 9600 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP
              16 4800 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP
              8 2400 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP
              4 1200 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 /n BlBP
              2 600 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP
              1 300 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP
              0.5 150 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP
              0.25 75 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP
              0.125 37.5 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP
              0.0625 18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
              0.0325 9.375 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP 0.122 / nBlBP 0.037 / nBlBP
              Subband bandwidth and spectral resolution options. Note that the table entries refer to the spacing between spectral channels -- that spacing is before any frequency smoothing, so these channels are notindependent.
              • nBlBP is the number of Baseline Board Pairs assigned to the subband.
              • Each subband may have a different number of spectral channels and polarization products, and each may be tuned independently.
              • There can be at most 16 subbands per baseband, and nBlBP must be an integer: 1, 2, 3, 4, 5, ..., 64.
              • The sum of nBlBP over all subbands must be less than or equal to 64.
              • Use of more than one BlBP for a subband may further restrict the number of subbands available in one or more of the basebands; see text for details.

              Here are four examples of allowed general observing setups which use all 64 BlBPs to produce the maximum number of channels times polarization products:

              Baseband Subband

              Pol'n

              Products

              Spectral

              channels

              nBlBP
              Example 1 A0/C0 sb0 RR 16384 64
              Example 2 A0/C0 sb0 RR 8192 32
              A0/C0 sb1 RR, LL 1024 8
              A0/C0 sb2 RR, LL 512 4
              B0/D0 sb0 RR, LL 2048 16
              B0/D0 sb1 RR,RL,LR,LL 256 4
              Example 3 A0/C0 sb0 RR 8192 32
              A0/C0 sb1 LL 1024 4
              A0/C0 sb2 RR, LL 1024 8
              A0/C0 sb3 RR,RL,LR,LL 1024 16
              A0/C0 sb4 RR,RL,LR,LL 256 4
              Example 4 A0/C0 sb0-5 RR,RL,LR,LL 64 6 x 1
              A0/C0 sb6 RR, LL 3840 1 x 30
              A0/C0 sb7 RR 768 1 x 3
              A0/C0 sb8 RR,RL,LR,LL 192 1 x 3
              B0/D0 sb0-2 RR,RL,LR,LL 64 3 x 1
              B0/D0 sb3 LL 768 1 x 3
              B0/D0 sb4 RR, LL 2048 1 x 16

               

              Recirculation

              Recirculation is a term to describe the method of increasing the number of spectral channels in a subband using correlator software (as opposed to Baseline Board stacking which uses correlator hardware, see below). Currently recirculation is achieved by limiting the subband bandwidth and thus only available for subbands less than 128 MHz wide.  When limiting the bandwidth in a subband, the correlator software can be directed to use the remaining CPU cycles on a Baseline Board pair to obtain more lags (in factors of two), running the data through the board for a second, third, etc., time; hence Recirculation.

              At some time in the future an alternative method of recirculation, using the extra CPU cycles freed up by increasing the integration time, would be made available. Recirculation by limiting of the subband bandwidth to increase the number of channels (in factors of two) was used in the pre-expansion VLA correlator.

              Recirculation versus Baseline Board stacking

              When faced with the choice between Recirculation and Baseline Board stacking (see below) to increase the number of channels in a subband for subbands narrower than 128 MHz we recommend the former, which is supported in observatory software (GOST, OPT). For 128 MHz subbands Baseline Board Stacking should be utilized to increase the number of channels.

              The current implementation of Recirculation is that for each halving of the subband bandwidth the number of channels in the subband may be doubled without having to trade off the use of other subbands. Because recirculation is achieved by limiting the subband bandwidth, it is not supported for 128 MHz subbands, whereas for 64 MHz subbands only a factor 2 recirculation is supported, etc. The maximum recirculation factor for a subband is 128/(subband bandwidth in MHz), and of course also subject to other configuration restrictions such as data rate.

              The juggling between the requested number of channels, subband bandwidth and the available number of Baseline Board pairs is dependent on the science goals and not easily formulated in a standard answer. However, if subbands of less than 128 MHz are used, Recirculation becomes an option for setups that can also be achieved with Baseline Board stacking. In such cases we suggest to use Recirculation where possible, and within the General Observing or Shared Risk requirements.  This frees up unused Baseline Board pairs for other use; alternatively, one becomes less dependent on all Baseline Board pairs being in working order.

              Recirculation with factors 8 to 64 is designated Shared Risk, and recirculation with factors over 64 Resident Shared Risk. The latter choice may have severe implications for the sensitivity as visibility integration time is used as trade-off. Ask the NRAO Helpdesk for more details.

              Baseline Board stacking

              As opposed to recirculation, which increases the number of channels in a subband by exploiting otherwise unused CPU resources, Baseline Board Stacking adds more channels to a subband by adding correlator hardware resources, i.e. using up more Baseline Board pairs.  Using Baseline Board stacking may therefore limit the number of subbands available in one or more of the basebands.  Understanding how this works requires understanding some of the details of the correlator hardware. That understanding is built into the GOST, and observers may simply use that tool to find out whether their particular setup will, in fact, work. But the results can be confusing without some understanding of the hardware constraints from which they arise.  These hardware constraints are complex, and most observers will not need to understand these details.  The following section is for those few who are attempting complex line experiments, and who find the GOST or the RCT restricting the number of subbands and/or channels they can use in unexpected ways  Most observers can skip it.

              Baseline Board Stacking and Correlator Use

              First let us consider how the correlator hardware is organized. The cross-multiplications in the WIDAR correlator are spread across 64 Baseline Board (BlB) pairs, arranged into 4 quadrants of 16 BlB pairs each. Each baseband is connected directly to one of those quadrants. In the simplest mode, each of the 16 BlB pairs of a quadrant handles the correlations for one of the 16 subbands of the corresponding baseband. Four basebands and four quadrants are required, in order to handle the full 8 GHz of bandwidth per polarization provided by the 3-bit (wideband) samplers: that 8 GHz is split into four basebands of 2 GHz each, with each baseband fed into a different BlB quadrant. Each BlB pair in that quadrant handles a subband of maximum bandwidth 128 MHz, so 16 BlB pairs handles 16 subbands for a total of 16x128 MHz= 2048 MHz.

              A single BlB pair produces 256 cross-correlations per baseline for a single subband, which can be used for a single polarization product (e.g., RR with 256 spectral channels), or two (RR+LL with 128 spectral channels each), or four (RR,RL,LR,LL, with 64 spectral channels each).

               

              When using the 8-bit samplers, the total bandwidth is only 2 GHz per polarization, split into two basebands of 1 GHz each. The simplest continuum setup uses only two quadrants, since there are only two basebands; and only 8 subbands are required to span the 8x128 MHz= 1024 MHz of each baseband. Three-quarters of the correlator BlB hardware remain unused.

               

               

              The spectral line mode allows access to these `extra' correlator resources through Baseline Board stacking: using multiple BlB pairs to process the same subband and produce more cross-correlations for that subband. This is done using crossbar switches which make the data for a single subband available to several BlB pairs. Those BlB pairs can then be used to produce more spectral channels for that subband, with n BlB pairs producing 256*n cross-correlations per baseline. The limit on the total number of cross-correlations (16384) stems from the total number of BlB pairs (64): 64x256= 16384.

              Unfortunately completely flexible crossbar switches are expensive, and could not be implemented in the VLA's new correlator. This means that one cannot route a given subband to a randomly-chosen BlB pair. The routings which are possible, are as follows:

              1. A subband in a baseband can be routed to any BlB pair within the corresponding quadrant.
              2. Data coming into a given BlB pair in one quadrant, can be routed to the corresponding BlB pair in any other quadrant.

               

              Routing option #1 means that one could use all the BlB pairs within a quadrant to correlate a single subband, yielding 16x256= 4096 cross-correlations for that subband:

               

              Routing option #2 means that one could use the BlB pairs in all 4 quadrants to correlate a single subband. One simple case would use 4 BlB pairs to correlate each of the 16 subbands in a single baseband, yielding 4x256= 1024 cross-correlations for each of those subbands. Note that in this case, no BlB pairs are left to correlate any data from the second baseband.

              Using routing option #2 does come with a subtle cost: assigning a BlB pair in quadrant X to correlate a subband corresponding to quadrant Y, removes that BlB pair from use in the baseband corresponding to quadrant X...and hence also removes the corresponding subband in that baseband. So getting more channels for a subband in one baseband, may prevent the use of a subband in a different baseband. To take a simple example, consider an experiment where one wishes to observe a single line in dual polarization with 512 channels (requiring 4 BlB pairs), plus as much continuum bandwidth as possible. Naively one would say there are 16 subbands in each baseband; one is used for the spectral line, so that leaves 16+15= 31 subbands, and with the widest subband bandwidth (128 MHz) the total available continuum should be 31x128 MHz= 3968 MHz per polarization. Actually however there are only 15+15 subbands available, or 30x128 MHz= 3840 MHz per polarization, because the spectral line subband has "eaten" one BlB pair corresponding to the other baseband:

               

              If the same spectral line required twice as many channels, this result in the loss of two subbands in both of the basebands:

              In some cases one may want to use a different routing, to use up subbands in one baseband in preference to another. For instance, the same spectral line setup (2048 cross-correlations for a single spectral line subband, plus as much continuum as possible) could be set up to allow 13 "continuum" subbands in the A0/C0 baseband, and the full 16 "continuum" subbands in B0/D0:

              Understanding these confusing constraints can help observers set up the VLA more effectively to achieve their scientific goals. For instance, for a mixed line+continuum experiment it works best to use the resource tools to set up the baseband tunings and subband channelization for the most important lines first, then add the desired continuum, and then see what correlator resources remain for any lines of secondary interest.

              The above examples all use BlB pair stacking in powers of 2, but this is not required.  To give some idea of more complex possibilities, the following tables give two examples of other possible configurations. The RCT display shows how the Baseline Boards are used to process the individual subbands. The cyan boxes show the Baseline Boards used to process data from baseband A0/C0, while the yellow boxes show Baseline Boards used to process data from baseband B0/D0.

              Complex Configuration Example #1
              BasebandSubbandPol'n productsSpectral channelsnBlBP
              A0/C0 sb0 RR 10240 40
              A0/C0 sb1 LL 768 3
              A0/C0 sb2 RR,LL 2176 17
              B0/D0 sb0 RR 256 1
              B0/D0 sb1 RR,LL 384 3
              RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

               

               

               

              Complex Configuration Example #2
              BasebandSubbandPol'n productsSpectral channelsnBlBP
              A0/C0 sb0 RR 4352 17
              A0/C0 sb1 RR, LL 1152 9
              B0/D0 sb0 RR,RL,LR,LL 192 3
              B0/D0 sb1 RR, LL 4480 35
              RCT display: corr-cfg-fig:sro2_8bit_ac17+9_bd3+35

               

              Note that the individual subbands can have different bandwidths, and those bandwidths may be chosen completely independently of the number of spectral channels in each subband. So for instance a subband with bandwidth 2 MHz and 1152 spectral channels would have a channel separation of 2 MHz/1152= 1.736 kHz; but the observer could equally well choose a bandwidth of 64 MHz for that subband, leading  to a channel separation of 64 MHz/1152=  55.56 kHz.

              Use of the 3-bit samplers further extends the possibilities.  Here is one example:

               

              3-bit Complex Configuration Example #1
              BasebandSubbandPol'n productsSpectral channelsnBlBP
              A1/C1 sb0-8 RR, LL, RL, LR 9 x 64 9 x 1
              A1/C1 sb9 RR, LL 1 x 1152 1 x 9
              A1/C1 sb10 RR 1 x 1792 1 x 7
              A1/C1 sb11 RR, LL 1 x 384 1 x 3
              A2/C2 sb0-11 RR, LL, RL, LR 12 x 64 12 x 1
              A2/C2 sb12 LL 1 x 768 1 x 3
              B1/D1 sb0-3 RR, LL, RL, LR 4 x 64 4 x 1
              B1/D1 sb4 RR, LL, RL, LR 1 x 320 1 x 5
              B2/D2 sb0-6 RR, LL, RL, LR 7 x 64 7 x 1
              B2/D2 sb7 RR, LL 1 x 640 1 x 5
              RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

               

              Once again, the GOST implements all of these constraints, and is generally smart enough to figure out the routing scheme that works best for your particular request.

              VLA Observational Status Summary 2016B

              VLA capabilities September 2016 - January 2017

              Introduction

              Purpose of Document, Older Versions of the OSS

              This Observational Status Summary (OSS) summarizes the instrumental status of the Karl G. Jansky Very Large Array (VLA) for the A configuration for the observing period 23 September 2016 through 23 January 2017 (Semester 2016B), and should be used when preparing proposals for the 1 February 2016 deadline. For capabilities prior to that date, we refer to our overview of all OSS versions available online.

              The OSS is intended as a ready reference for those contemplating use of the VLA for their astronomical research. The information is in summary form; those requiring greater detail should use the NRAO Helpdesk, or refer to the manuals and documentation listed in Documentation. Most of the information contained here, and much more, is available through the VLA science web page and the companion VLBA science web page.

              The VLA is a large and complex modern instrument. Some familiarity with the principles and practices of its operation is necessary for its efficient use. Although the NRAO strives to make using the VLA as simple as possible, users must be aware that proper selection of observing mode and calibration technique is often crucial to the success of an observing program. Inexperienced and first–time users are encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing program. For more details, refer to the Visiting the DSOC and VLA page. The VLA is an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it.

               

              Overview of the VLA

              The Karl G. Jansky Very Large Array (VLA) is a 27–element interferometric array, arranged along the arms of an upside-down Y, which produces images of the radio sky at a wide range of frequencies and resolutions. The VLA is located at an elevation of 2100 meters on the Plains of San Agustin in southwestern New Mexico, and is managed from the Pete V. Domenici Science Operations Center (DSOC) in Socorro, New Mexico.

              The basic data produced by the VLA are the visibilities, or measures of the spatial coherence function, formed by correlation of signals from the array's elements. The most common mode of operation will use these data, suitably calibrated, to form images of the radio sky as a function of sky position and frequency. Another mode of observing, commonly called phased array, allows operation of the array as a single element through coherent summation of the individual antenna signals. This mode is most commonly used for Very Long Baseline Interferometry (VLBI) observing and for observations of rapidly varying objects, such as pulsars.

              The VLA can vary its resolution over a range exceeding a factor of ∼50 through movement of its component antennas. There are four basic antenna arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. These configurations are denoted D, C, B, and A, respectively. For details about antenna positions in the various configurations we refer to the stations position file (pdf version).

              The VLA completes one cycle through all four configurations in an approximately 16 month period. Consult the Configuration Plans and Proposal Deadlines page or recent NRAO and AAS newsletters for current and up-to-date configuration schedules and associated proposal deadlines. Refer to the Guide to Proposing for the VLA for information on how to submit an observing proposal.

              Observing projects on the VLA will vary in duration from as short as 1/2 hour to as long as several weeks. Most observing runs have durations of a few to 24 hours with only one, or perhaps a few, target sources. However, since the VLA is a two-dimensional array, images can be made with data durations of less than one minute. This mode, commonly called snapshot mode, is well suited to surveys of relatively strong, isolated objects. See the section on Snapshots for more detail.

              All VLA antennas are outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers cover the frequency ranges of 1–2 GHz (L-band), 2–4 GHz (S-band), 4–8 GHz (C-band), 8–12 GHz (X-band), 12–18 GHz (Ku-band), 18–26.5 GHz (K-band), 26.5–40 GHz (Ka-band), and 40–50 GHz (Q-band). Additionally, all antennas of the VLA have receivers for lower frequencies, enabling observations at P-band (230–470 MHz). These low frequency receivers also work at 4-band (54–86 MHz), and new feeds have been deployed on a small number of VLA antennas to observe at this frequency range.

              The VLA correlator is both powerful and flexible. Details of the correlator configurations being offered for VLA science are described in the WIDAR Section of this document. It is important to realize that the VLA correlator is fundamentally a spectral line correlator and that even continuum observations are done in a wide-band mode with many channels.

              Proposing for the VLA

              The Call for Proposals

              The most recent Call for Proposals details the General Observing (GO) capabilities being offered for the Karl G. Jansky Very Large Array (VLA).

              In addition to these general capabilities, NRAO continues to offer shared risk observing options for those who would like to push the capabilities of the VLA beyond those offered for general use. These are the Shared Risk Observing (SRO) and Resident Shared Risk Observing (RSRO) programs.

              Details about what is being offered for each program is given below. If you have any questions or problems with any link or tool, please submit a ticket through the NRAO Helpdesk.

              Considering the lack of future hybrid configurations after semester 2016A, guidelines on how to substitute such configurations with the use of principal array configurations are presented in the Array Configurations section of the Guide to Proposing for the VLA.


              General Observing (GO) and Shared-Risk Observing (SRO)

              Summary of Capabilities

              As described in the Call for Proposals, the VLA offers continuous frequency coverage from 1–50 GHz in the following observing bands: 1–2 GHz (L-band); 2–4 GHz (S-band); 4–8 GHz (C-band); 8–12 GHz (X-band); 12–18 GHz (Ku-band); 18–26.5 GHz (K-band); 26.5–40 GHz (Ka-band); and 40–50 GHz (Q-band). Both single pointing and mosaics with discrete, multiple field centers will be supported under General Observing (GO). P-band (230–470 MHz) Stokes I continuum observations are also offered under GO; see the Low Frequency Observing section below for more details. Data rates of up to 60 MB/s (216 GB/hour) will be available to all users, combined with correlator integration time limits per band and per configuration, as described in the Time Resolution and Data Rates section. However, data rates in excess of 25 MB/s (90 GB/hour) require additional justification. Limitations on frequency settings and tuning ranges are described in the Frequency Bands and Tunability section.

              The GO capabilities being offered are:

              Capability Description
              8-bit samplers
              • Standard default setups for:
                • 2 GHz bandwidth continuum observations at S/C/X/Ku/K/Ka/Q bands (16 × 128 MHz subbands)
                • 1 GHz bandwidth continuum observations at L band (16 × 64 MHz subbands)
                • 256 MHz bandwidth continuum observations at P band (16 × 16 MHz subbands)
              • Flexible setups for spectroscopy using two, independently tunable 1 GHz baseband pairs, each of which can be split into up to 16, flexibly tunable, subbands
              • Single, dual, and full polarization products
              3-bit samplers
              • Standard default setups for:
                • 8 GHz bandwidth continuum observations at K/Ka/Q bands
                • 6 GHz bandwidth at Ku band
                • 4 GHz bandwidth at C/X bands
              • Flexible setups for spectroscopy using four, independently tunable 2 GHz baseband pairs, each of which can be split into up to 16, flexibly tunable, subbands
              • Single, dual, and full polarization products
              Mixed 3-bit and 8-bit samplers
              • Allows more flexibility for simultaneous continuum and high-resolution spectral line observing

              Subarrays

              • Up to 3 independent subarrays using standard 8-bit continuum setups
              Phased-array for VLBI
              • VLA Phased Array (Y27) for VLBI (see the VLBA call for proposals and the VLBA Observational Status Summary for a detailed description of the VLA Phased Array capabilities being offered)


              SRO capabilities can be set up via the Observing Preparation Tool (OPT) and run through the dynamic scheduler without intervention, but are not as well tested as GO capabilities. A summary of the SRO capabilities being offered are:

              • On-the-Fly (OTF) mosaicing
              • Up to 32 subbands per baseband with the 8-bit samplers

              ‡Note: OTF mosaic observations have certain limitations on the length and number of scans. Contact the VLA staff through the NRAO Helpdesk to check if the OTF mosaic observations to be proposed are technically feasible ahead of submitting a proposal.

              We expect that most SRO programs will have no or only minor problems that can be corrected quickly. If an SRO program fails, however, and it becomes clear that detailed testing with additional expertise is needed, then the project must make an experienced member from their team available to help troubleshoot the problem. In some cases, this may require the presence of that experienced member in Socorro. If adequate support from the project is not given, then the time on the telescope will be forfeited. The additional effort is to be determined based on discussions with the NRAO staff and management and the project team.

              The guidelines for General and Shared Risk observing proposals, along with information about tools and other advice, can be found in the VLA Proposal Submission Guidelines.

              Resident Shared Risk Observing (RSRO)

              Summary of Capabilities

              The RSRO program provides access to extended capabilities of the VLA that require additional testing in exchange for a period of residence to help commission those capabilities. Capabilities that would fall under the RSRO program include:

              • correlator dump times shorter than 50 msec, including integration times as short as 5 msec for transient detection;
              • pulsar observations;
              • data rates above 60 MB/s;
              • recirculation beyond a factor of 64 in the correlator;
              • P-band system (230–470 MHz) polarimetry and spectroscopy;
              • 4-band system (54–86 MHz, see the Low Frequency Observing section below);
              • more than 3 subarrays or subarrays with the 3-bit system;
              • complex phased array observations (e.g., pulsar and complex VLBI observing modes), and;
              • frequency averaging in the correlator: a new capability for averaging to wider frequency channels in the correlator has been developed that will reduce the data volume for all continuum subbands

               

              The guidelines for Resident Shared Risk Observing proposing, along with requirements and considerations, can be found in the VLA Proposal Submission Guidelines.

               

              Low Frequency Observing (P-band and 4-band) and VLITE

              The new, low frequency receiver system, developed in collaboration with the Naval Research Laboratory, will be available for Stokes I continuum observations at P-band (230–470 MHz) through the GO program. Use of the P-band system for polarimetry and/or spectroscopy will be through the RSRO program as noted above. The new receivers also work at 4-band (54–86 MHz), and new feeds will be deployed on at least ten (and up to fourteen) VLA antennas. Both 4-band and P-band can be observed simultaneously, but access to the 4-band system is only available through the RSRO program.

              Finally, the commensal VLITE system will take data at P-band during regular observations that use bands other than P-band. The VLITE system is deployed on ten VLA antennas. Observers wishing to gain access to the commensal VLITE data taken during their VLA observations should follow the instructions on the VLITE web page for doing so.

               

              To report errors or problems encountered in any link or while using any NRAO tool listed here, please submit a ticket through the NRAO Helpdesk.

              Performance of the VLA during the Next Semester

              Resolution

              Resolution

              The VLA's resolution is generally diffraction-limited, and thus is set by the array configuration and frequency of observation. It is important to be aware that a synthesis array is blind to structures on angular scales both smaller and larger than the range of fringe spacings given by the antenna distribution. For the former limitation, the VLA acts like any single antenna—structures smaller than the diffraction limit (θ ∼ λ/Bmax) are not seen—with the resulting image smoothed to the resolution of the array. The latter limitation is unique to interferometers; it means that structures on angular scales significantly larger than the fringe spacing formed by the shortest baseline are not measured. No subsequent processing can fully recover this missing information, which can only be obtained by observing in a more compact array configuration, by using the mosaicking method, or by utilizing data from an instrument such as a large single antenna or a compact array comprising smaller antennas which provides this information.

              Table 3.1.1 summarizes the relevant resolution information. This table shows the maximum and minimum antenna separations, the approximate synthesized beam size (full width at half-power) for the central frequency for each band, and the scale at which severe attenuation of large scale structure occurs.

              Table 3.1.1: Configuration Properties
              ConfigurationABCD
              Bmax (km1) 36.4 11.1 3.4 1.03
              Bmin (km1) 0.68 0.21 0.0355 0.035
              Band Synthesized Beamwidth θHPBW(arcsec)1,2,3
              74 MHz (4) 24 80 260 850
              350 MHz (P) 5.6 18.5 60 200
              1.5 GHz (L) 1.3 4.3 14 46
              3.0 GHz (S) 0.65 2.1 7.0 23
              6.0 GHz (C) 0.33 1.0 3.5 12
              10 GHz (X) 0.20 0.60 2.1 7.2
              15 GHz (Ku) 0.13 0.42 1.4 4.6
              22 GHz (K) 0.089 0.28 0.95 3.1
              33 GHz (Ka) 0.059 0.19 0.63 2.1
              45 GHz (Q) 0.043 0.14 0.47 1.5
              Band Largest Angular Scale θLAS(arcsec)1,4
              74 MHz (4) 800 2200 20000 20000
              350 MHz (P) 155 515 4150 4150
              1.5 GHz (L) 36 120 970 970
              3.0 GHz (S) 18 58 490 490
              6.0 GHz (C) 8.9 29 240 240
              10 GHz (X) 5.3 17 145 145
              15 GHz (Ku) 3.6 12 97 97
              22 GHz (K) 2.4 7.9 66 66
              33 GHz (Ka) 1.6 5.3 44 44
              45 GHz (Q) 1.2 3.9 32 32
              These estimates of the synthesized beamwidth are for a uniformly weighted, untapered map produced from a full 12 hour synthesis observation of a source which passes near the zenith.
              Notes:
              1. Bmax is the maximum antenna separation, Bmin is the minimum antenna separation, θHPBW is the synthesized beam width (FWHM), and θLAS is the largest scale structure visible to the array.
              2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees.
              3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θHPBW. The values for snapshots are about 1.3 times the listed values.
              4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations, the quoted numbers should be divided by two.
              5. For the C configuration, an antenna from the middle of the north arm is moved to the central pad N1. This results in improved imaging for extended objects, but may slightly degrade snapshot performance. Note that although the minimum spacing is the same as in D configuration, the surface brightness sensitivity and image fidelity to extended structure is considerably inferior to that of the D configuration.

              The following figure is a graphical representation of the synthesized beamwidths for natural and robust weighting for the four main array configurations between 1 and 50 GHz. Also available are synthesized beamwidth figures for the low frequency (1–12 GHz) and the high frequency (12–50 GHz) receiver bands.

              A project with the goal of doubling the longest baseline available in the A configuration, by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town, was established in the late 1990s and used through 2005. This link is no longer operational; the objective of implementing a new digital Pie Town link, now that EVLA construction is complete, remains unfunded.

              Sensitivity

              Sensitivity

              The theoretical thermal noise expected for an image using natural weighting of the visibility data is given by:

              where:

              - SEFD is the system equivalent flux density (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEFD indicate more sensitive performance. For the VLA's 25–meter paraboloids, the SEFD is given by the equation SEFD = 5.62TsysA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band.
              - ηc is the correlator efficiency (~0.93 with the use of the 8-bit samplers).
              - npol is the number of polarization products included in the image; npol = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in RCP or LCP.
              - N is the number of antennas.
              - tint is the total on-source integration time in seconds.
              - Δν is the bandwidth in Hz.

              Figure 3.2.1 shows the SEFDs as a function of frequency used in the VLA exposure calculator for those Cassegrain bands currently installed on VLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Figure 3.2.2 shows the SEFDs as a function of frequency for the P-band. These measurements are based on imaging of a field far from the galactic plane. Table 3.2.1 gives the SEFD at some fiducial VLA frequencies.

              Figure 3.2.1: SEFD used in the Exposure Calculator for the VLA. Above left: The system equivalent flux density as a function of frequency for the L, S, C and X-band receivers. Right: The system equivalent flux density as a function of frequency for the Ku, K, Ka, and Q-band receivers. SEFDs at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere and were determined under good conditions. 

              Figure 3.2.2: The SEFD used in the VLA Exposure Calculator as a function of frequency for the P-band receiver

              Note that the theoretical rms noise calculated using equation 1 is the best limit possible. There are several factors that will tend to increase the noise compared with theoretical:

              • For the more commonly used robust weighting scheme, intermediate between pure natural and pure uniform weightings (available in the AIPS task IMAGR and CASA task clean), typical parameters will result in the sensitivity being a factor of about 1.2 worse than the listed values.
              • Confusion. There are two types of confusion: (i) that due to confusing sources within the synthesized beam, which affects low resolution observations the most. Table 3.2.1 shows the confusion noise in D configuration (see Condon 2002, ASP Conf. 278, 155), which should be added in quadrature to the thermal noise in estimating expected sensitivities. The confusion limits in C configuration are approximately a factor of 10 less than those in Table 3.2.1; (ii) confusion from the sidelobes of uncleaned sources lying outside the image, often from sources in the sidelobes of the primary beam. This confusion primarily affects low frequency observations.
              • Weather. The sky and ground temperature contributions to the total system temperature increase with decreasing elevation. This effect is very strong at high frequencies, but is relatively unimportant at the other bands. The extra noise comes directly from atmospheric emission: primarily from water vapor at K-band, and from water vapor and the broad wings of the strong 60 GHz O2 transitions at Q-band.
              • Losses from the 3-bit samplers. The VLA's 3-bit samplers incur an additional 10–15% loss in sensitivity above that expected—i.e., the efficiency factor ηc = 0.78 to 0.83.

               

              Table 3.2.1: SEFDs and D-Configuration Confusion Limits
              Frequency SEFD
              (Jy)
               RMS confusion level
              in D config (µJy/beam)
              0.39 GHz (P) 2790 4200
              1.5 GHz (L) 420 89
              3.0 GHz (S) 370 14
              6.0 GHz (C) 310 2.3
              10.0 GHz (X) 250 negligible
              15 GHz (Ku) 350 negligible
              22 GHz (K) 560 negligible
              33 GHz (Ka) 710 negligible
              45 GHz (Q) 1260 negligible

               

              In general, the zenith atmospheric opacity to microwave radiation is very low: typically less than 0.01 at L, C, and X-bands; 0.05 to 0.2 at K-band; and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Observing conditions are best at night and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable.

              Observers should remember that clouds, especially clouds with large water droplets (thunderstorms), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz.

              Tipping scans—which are currently unavailable but will be implemented at some time in the future—can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the complex gain calibrator (amplitude and phase) and target source are observed.

              When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following items:

              • Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target.

                Both CASA and AIPS allow the application of elevation-dependent gains and an estimated opacity generated from ground-based weather through the CASA tasks gencal and plotweather, and AIPS task INDXR.

              • Pointing. The SEFD quoted above assumes good pointing. Under calm, nighttime conditions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime can be much worse—occasionally exceeding 1 arcminte due to the effects of solar heating of the antenna structures. Moderate winds have a very strong effect on both pointing and antenna figure. The maximum wind speed recommended for high frequency observing is 11 mph (5 m/s). Wind speeds near the stow limit 45 mph (20 m/s) will have a similar negative effect at 8 and 15 GHz.

              To achieve increased pointing accuracy, referenced pointing is recommended where a nearby calibrator is observed in interferometric pointing mode every hour or so. The local pointing corrections measured can then be applied to subsequent target observations. This reduces rms pointing errors to as little as 2–3 arcseconds (but more typically 5–7 arcseconds) if the reference source is within about 15 degrees in azimuth and elevation of the target source and the source elevation is less than 70 degrees. At source elevations greater than 80 degrees (zenith angle < 10 degrees), source tracking becomes difficult; it is recommended to avoid such source elevations during the observation preparation setup.

              Use of referenced pointing is highly recommended for all Ku, K, Ka, and Q-band observations, and for lower frequency observations of objects whose total extent is a significant fraction of the antenna primary beam. It is usually recommended that the referenced pointing measurement be made at 8 GHz (X-band), regardless of what band your target observing is at, since X-band is the most sensitive and the closest calibrator is likely to be weak. Proximity of the reference calibrator to the target source is of paramount importance; ideally the pointing sources should precede the target by 20 or 30 minutes in Right Ascension (RA). The calibrator should have at least 0.3 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution.

              To aid VLA proposers there is an online guide to the exposure calculator; the exposure calculator provides a graphical user interface to these equations.

              Special caveats apply for P-band (230–470 MHz) observing. The SEFD's in Figure 3.2.2 or that listed in table 3.1.2 are from an observation taken far from the Galactic plane, where the sky brightness is about 30K. At P-band, Galactic synchrotron emission is very bright in directions near the Galactic plane. The system temperature increase due to Galactic emission will degrade sensitivity by factors of two to three for observations in the plane, and by a factor of five or more at or near the Galactic center. Additionally, the antenna efficiency (currently about 0.31 for 300 MHz) will decline with both increasing and decreasing frequencies from the center of P-band.

              The beam-averaged brightness temperature measured by a given array depends on the synthesized beam, and is related to the flux density per beam by:

              where Tb is the brightness temperature (Kelvins) and Ω is the beam solid angle. For natural weighting (where the angular size of the approximately Gaussian beam is ∼1.5λ/Bmax), and S in mJy per beam, the parameter F depends on the synthesized beam, therefore on the array configuration, and has the approximate value of F = 190, 18, 1.7, 0.16 for A, B, C, and D configurations, respectively. The brightness temperature sensitivity can be obtained by substituting the rms noise, ΔIm, for S. Note that Equation 2 is a beam-averaged surface brightness; if a source size can be measured, then the source size and integrated flux density should be used in Equation 2 and the appropriate value of F calculated. In general, the surface brightness sensitivity is also a function of the source structure and how much emission may be filtered out due to the sampling of the interferometer. A more detailed description of the relation between flux density and surface brightness is given in Chapter 6 of Reference 1, listed in Documentation.

              For observers interested in HI in galaxies, a number of interest is the sensitivity of the observation to the HI mass. This is given by van Gorkom et al. (1986; AJ, 91, 791):

              where D is the distance to the galaxy in Mpc, and SΔV is the HI line area in units of Jy km/s.

               

              VLA Frequency Bands and Tunability

              Bands

              For observations taken with the 8-bit samplers, each receiver can tune to two different frequencies, each 1024 MHz wide, within the same frequency band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies, except for the low-band receiver (50–500 MHz), which provides linear polarization (X and Y). Each of these four data streams follows the VLA nomenclature and are known as IF (for Intermediate Frequency channel) A, B, C, and D. IFs A and B provide RCP (or Y when applicable), IFs C and D provide LCP (or X when applicable). IFs A and C are always at the same frequency, as are IFs B and D (but note that the A and C IFs frequency is usually different from the B and D frequency). We normally refer to these two independent data streams as IF pairs, i.e., the A/C pair and the B/D pair. In 8-bit mode, a maximum of 1024 MHz can be correlated for each IF pair (see the WIDAR Section), for a total maximum bandwidth of 2048 MHz. To distinguish this 8-bit system from the 3-bit system, these IF pairs are denoted A0/C0 and B0/D0.

              More options are available with the 3-bit samplers. This system provides four (R,L) polarization pairs, each 2048 MHz wide. The A/C IF pair provides two sampled pairs, labelled A1/C1 and A2/C2, and the B/D IF pair provides two sampled pairs, labelled B1/D1 and B2/D2.

              For more details on the 8-bit and 3-bit samplers see the VLA Samplers section.

              The tuning ranges, along with default frequencies for continuum applications, are given in Table 3.3.1 below.

              Table 3.3.1: Default frequencies for continuum applications
              BandRange18-bit continuum applications (GHz)3-bit continuum applications (GHz)

              (GHz)IF pair A0/C0IF pair B0/D0IF pair A1/C1IF pair A2/C2IF pair B1/D1IF pair B2/D2
              4 m (4) 0.058 – 0.0842 .054 – .086
              90 cm (P) 0.22 – 0.503 0.224 – 0.4803
              20 cm (L) 1.0 – 2.04 1.0 – 1.54 1.5 – 2.04
              13 cm (S) 2.0 – 4.0 2.0 – 3.0 3.0 – 4.0
              6 cm (C) 4.0 – 8.0 4.5 – 5.5 5.5 – 6.5 4.0 – 6.0 6.0 – 8.0
              3 cm (X) 8.0 – 12.0 8.0 – 9.0 9.0 – 10.0 8.0 – 10.0 10.0 – 12.0
              2 cm (Ku) 12.0 – 18.0 13.0 – 14.0 14.0 – 15.0 12.0 – 14.0 14.0 – 16.0 16.0 – 18.0
              1.3 cm (K) 18.0 – 26.5 20.2 – 21.2 21.2 – 22.2 22.0 – 24.0 24.0 – 26.0 18.0 – 20.0 20.0 – 22.0
              1 cm (Ka) 26.5 – 40.0 32.0 – 33.0 31.0 – 32.0 33.0 – 35.0 35.0 – 37.0 29.0 – 31.0 31.0 – 33.0
              0.7 cm (Q) 40.0 – 50.0 40.0 – 41.0 41.0 – 42.0 44.0 – 46.0 46.0 – 48.0 40.0 – 42.0 42.0 – 44.0

              Notes:

              1.  Listed here are the nominal band edges. For all bands, the receivers can be tuned to frequencies outside this range, but at the cost of diminished performance. Contact the NRAO Helpdesk for further information.
              2. The 4-band system is currently under development. Observing time may be requested through the RSRO program.
              3. The default setup for P-band will provide 16 subbands from the A0/C0 IF pair, each 16 MHz wide, to cover the frequency range 224–480 MHz. The channel resolution is 125 kHz. 
              4. The default frequency setup for L-band comprises two 512 MHz IF pairs (each comprising 8 contiguous subbands of 64 MHz) to cover the entire 1–2 GHz of the L-band receiver.

               

              Tuning Restrictions

              In general, for all frequency bands except Ka, if the total span of the two independent IF pairs of the 8-bit system (defined as the frequency difference between the lower edge of one IF pair and the upper edge of the other) is less than 8.0 GHz, there are no restrictions on the frequency placements of the two IF pairs. For K, Ka, and Q-bands—the only bands where a span greater than 8 GHz is possible—there are special rules:

              • At Ka-band, the low frequency edge of the A0/C0 IF pair must be greater than 32.0 GHz. There is no restriction on the B0/D0 frequency, unless the B0/D0 band overlaps the A0/C0 band when the latter is tuned at or near the 32.0 GHz limit. In this case, the Observation Preparation Tool (OPT) may not allow the requested frequency setups. Users wanting to use such a frequency setup are encouraged to contact the NRAO Helpdesk for possible tuning options.
              • At K and Q-bands, if the frequency span is greater than 8.0 GHz, the B0/D0 frequency must be lower than the A0/C0 frequency.

              For the 3-bit system, the maximum frequency span permitted for the A1/C1 and A2/C2 IF pairs is about 5000 MHz. The same restriction applies to B1/D1 and B2/D2. The tuning restrictions given above for the separation and location of the 8-bit pairs A0/C0 and B0/D0 also apply to the 3-bit pairs, with A0/C0 replaced by A1/C1 and A2/C2, and B0/D0 replaced by B1/D1 and B2/D2.

               

              VLA Samplers

              The VLA is equipped with two different types of samplers, 8-bit with 1GHz bandwidth, and 3-bit with 2GHz bandwidth. The choice depends on your science goals and on technicalities described below.

              The 8-bit Set consists of four 8-bit samplers running at 2.048 GSamp/sec. The four samplers are arranged in two pairs, each pair providing 1024 MHz bandwidth in both polarizations. The two pairs are denoted A0/C0 and B0/D0. Taken together, the four samplers offer a maximum of 2048 MHz coverage with full polarization. The frequency spans sampled by the two pairs need not be adjacent. Some restrictions apply, depending on band, as described in the section on Frequency Bands and Tunability.

              The 3-bit Set consists of eight 3-bit samplers running at 4.096 GSamp/sec. The eight samplers are arranged as four pairs, each pair providing 2048 MHz bandwidth in both polarizations. Two of these pairs, denoted A1/C1 and A2/C2 cannot span more than 5000 MHz (lower edge of one to the higher edge of the other). The same limitation applies to the second pair, denoted B1/D1 and B2/D2. The tuning restrictions are described in the section on Frequency Bands and Tunability.

               

              Which set to use?

              • S, L, and 4/P-band observations, whether line or continuum, should use the 8-bit sampler set.
              • C and X-band continuum observations should use 3-bit samplers in order to exploit the full 4 GHz bandwidth: in spite of the 15% reduction in sensitivity that comes with 3-bit (at equal bandwidth to the 8-bit samplers—see below for details) and the reduced effective bandwidth after removing RFI, this still provides superior overall sensitivity. For more details we refer to EVLA memo 166.
              • Ku, K, Ka, and Q-band continuum observations should use the 3-bit samplers for maximum bandwidth.
              • Wide-band spectral line searches requiring more than 2 GHz span should use the 3-bit samplers.
              • Spectral-line observations which fit within two, possibly disjoint, 1 GHz bands should use the 8-bit set.
              • Simultaneous continuum and high resolution spectral line observation can use mixed 3-bit and 8-bit samplers. The 3-bit samplers in this case will be set up to deliver the continuum data, while the 8-bit samplers will be for the spectral line data. This mix mode can be used in C-band and higher.


              Major Characteristics of each Set

              The 8-bit samplers are warranted for observations at 4/P, L, and S-bands. The full analog bandwidth from the receivers fits within the 2048 MHz span covered by the samplers.

              For the 3-bit samplers, users need to be aware of the following issues:

              • Sensitivity: compared to the 8-bit system, the sensitivity of the 3-bit samplers is worse by ~15% (at equal bandwidth). Alternatively, a given continuum noise level requiring on-source integration time T with the 8-bit (two bands of 1GHz), requires 0.33T with the 3-bit (4 bands of 2GHz, assuming the bandwidth is available from the front end).
              • Resonances: each of the eight 3-bit samplers on an antenna has a resonance about 3 MHz wide. Each resonance is independent of all others, so there is no correlated signal between antennas. The resonance degrades the spectrum in its narrow frequency range, but has little effect on continuum observing. Bandpass solutions will be affected, but can be interpolated over. Spectral-line calibration and images at the affected frequencies will show significant loss in sensitivity. The resonances are easily seen in autocorrelation spectra, and it is recommended that users, especially spectral-line users, utilize these to locate the compromised frequencies.
              • Amplitude Calibration: The traditional method for both 8- and 3-bit systems is to observe a flux-density calibrator, use self-cal to determine the antenna amplitude calibration factors (gains), and transfer the gains to the phase calibrator and target. For 3-bit samplers this procedure gives results good to 5% between elevations of 20–70 degrees. (Expect worse at the upper edge of Q-band and/or during bad weather.) The switched power data can be used to correct for system gain variations and works well for the 8-bit samplers. For 3-bit samplers, the Pdif depends on the Psum, i.e., Pdif is non-linear and its application will bias the resulting visibilities by 5–10%. The origin of this effect is understood, but we have not yet determined how best to compensate for it. Because of this, we do not recommend use of the Psum and Pdif data to calibrate visibilities from the 3-bit samplers. We do, however, recommend that the requantizer gains in the switched power data be applied to remove gain changes. For more information about the switched power, Psum, and Pdif, see EVLA memo 145.

               

              Setting up the 8-bit or 3-bit Samplers

              Either set requires an initial scan for each individual LO (frequency) tuning, during which power levels are optimized.

              For the 8-bit system, a dummy scan of 1 minute duration is sufficient for each tuning. This  is usually done while the antennas are slewing at the start of an observing file, as the pointing direction of the antennas is not critical.

              For the 3-bit system, the requirements are more demanding, see the section on 3-bit setup within the Guide to Observing with the VLA. The minimum setup time is 1 minute for each tuning to adjust the power levels and bandpass slopes across the 2GHz samplers. These values are retained and applied if the tuning is re-encountered in the same observation. Additionally, every time the LO setup is changed—whether or not it is new (e.g., changing from 8-bit X-band reference pointing back to target)—a scan of 30 seconds is needed to reset the subband gains (requantizers) in the correlator. For better amplitude calibration at high frequencies, the 3-bit initial setup should be near the elevation of the target, so do it after the first 8-bit setup described above. For 3-bit observing without 8-bit (e.g., C or X-band without reference pointing), the power variation with elevation is small, so the 3-bit setup can be done at any elevation.

              For settings that use a mix of 3-bit and 8-bit samplers, the guidelines to set up the 3-bit samplers should be followed.


              Other issues

              The overhead for setup of 3-bit samplers can eat into observing time, especially for projects with many different LO settings, and/or sources all over the sky accompanied by band change, reference pointing, and requantizer reset for each direction. The impact is most severe for short scheduling blocks.

              Polarization testing conducted so far indicates no degradation of performance by using the 3-bit samplers.

              Field of View

              Primary Beam

              The ultimate factor limiting the field of view is the diffraction-limited response of the individual antennas. An approximate formula for the full width at half power in arcminutes is: θPB = 45/νGHz. More precise measurements of the primary beam shape have been derived and are incorporated in AIPS task PBCOR and CASA task clean (and the imaging toolkit) to allow for correction of the primary beam attenuation in wide-field images. Objects larger than approximately half this angle cannot be directly observed by the array. However, a technique known as mosaicking, in which many different pointings are taken, can be used to construct images of larger fields. Refer to References 1 and 2 in Documentation for details.

              Guidelines for mosaicking with the VLA are given in the Guide to Observing with the VLA.

               

              Chromatic Aberration (Bandwidth Smearing)

              The principles upon which synthesis imaging are based are strictly valid only for monochromatic radiation. When visibilities from a finite bandwidth are gridded as if monochromatic, aberrations in the image will result. These take the form of radial smearing which worsens with increased distance from the delay-tracking center. The peak response to a point source simultaneously declines in a way that keeps the integrated flux density constant. The net effect is a radial degradation in the resolution and sensitivity of the array.

              These effects can be parameterized by the product of the fractional bandwidth (Δν/ν0) with the source offset in synthesized beamwidths (θ0HPBW). Table 3.5.1 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter and should be used to determine how much spectral averaging can be tolerated when imaging a particular field.

              Table 3.5.1: Reduction in Peak Response Due to Bandwidth Smearing
              (Δν/ν0)*(θ0HPBW) Peak Width
              0.0 1.00 1.00
              0.50 0.95 1.05
              0.75 0.90 1.11
              1.0 0.80 1.25
              2.0 0.50 2.00

              Note: The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). Δν/ν0 is the fractional bandwidth; θ0HPBW is the source offset from the phase tracking center in units of the synthesized beam.

               

              Time-Averaging Loss

              The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parametrized, except for observations at δ = 90°. In this case, the effects are identical to the bandwidth effect except they operate in the azimuthal, rather than the radial, direction. The functional dependence is the same as for chromatic aberration with Δν/ν0 replaced by ωeΔtint, where ωe is the Earth's angular rotation rate, and Δtint is the averaging interval.

              For other declinations, the effects are more complicated and approximate methods of analysis must be employed. Chapter 13 of Reference 1 in Documentation considers the average reduction in image amplitude due to finite time averaging. The results are summarized in Table 3.5.2, showing the time averaging in seconds which results in 1%, 5% and 10% loss in the amplitude of a point source located at the first null of the primary beam. These results can be extended to objects at other distances from the phase tracking center by noting that the loss in amplitude scales with (θΔtint)2, where θ is the distance from the phase center and Δtint is the averaging time. We recommend that observers reduce the effect of time-average smearing by using integration times as short as 1 or 2 seconds (also see the section on Time Resolution and Data Rates) in the A and B configurations.

              Table 3.5.2: Averaging Time for a Given Amplitude Loss
              Amplitude loss
              Configuration 1.0% 5.0% 10.0%
              A 2.1 4.8 6.7
              B 6.8 15.0 21.0
              C 21.0 48.0 67.0
              D 68.0 150.0 210.0

              Note: The averaging time (in seconds) results in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines.

               

              Non-Coplanar Baselines

              The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the VLA with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at 90 cm and 20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations.

              Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS task IMAGR and CASA task clean. Refer to the package help files for these tasks, or consult with the NRAO Helpdesk for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the W-projection method available in CASA (see EVLA Memo 67 for more details).

              Time Resolution and Data Rates

              The default integration times for the various array configurations and frequency bands are as follows:

              Table 3.6.1: Default Integration Times
              ConfigurationsObserving
              Bands
              Default
              integration time
              A, B, C, D P 2 seconds
              A L S C X Ku K Ka Q 2 seconds
              B L S C X Ku K Ka Q 3 seconds
              C, D X Ku K Ka Q 3 seconds
              C, D L S C 5 seconds

              Observations with the 3-bit (wideband) samplers, when applicable, must use these integration times. Observations with the 8-bit samplers may use shorter integration times, but these must be requested and justified explicitly in the proposal, and obey the following restrictions:

              Table 3.6.2: Minimum integration times and maximum data rates
              Proposal type

              Minimum integration time

              Maximum data rate
              General Observing (GO) and
              Shared Risk Observing (SRO)
              50 msec 25 MB/s (90 GB/hr), or up to
              60 MB/s (216 GB/hr) with additional justification
              Resident Shared Risk Observing (RSRO) < 50 msec > 60 MB/s (216 GB/hr)

              Note that integration times as short as 5 msec and data rates as high as 300 MB/s can be supported for some observing, though any such observing is considered Resident Shared Risk Observing (RSRO). For these short integration times and high data rates there will be limits on bandwidth and/or number of antennas involved in the observation. Those desiring to utilize such short integration times and high data rates should consult with NRAO staff.

              The maximum recommended integration time for any VLA observing is 60 seconds. For high frequency observations with short scans, e.g., fast switching (as described in Rapid Phase Calibration and the Atmospheric Phase Interferometer (API)), shorter integration times may be preferable.

              Observers should bear in mind the data rate of the VLA when planning their observations. For Nant antennas and integration time Δt, the data rate is:

              Data rate ~ 45 MB/sec × (Nchpol/16384) × Nant × (Nant − 1)/(27×26) / (Δt/1 sec)
              ~ 160 GB/hr × (Nchpol/16384) x Nant × (Nant − 1)/(27×26) / (Δt/1 sec)
              ~ 3.7 TB/day × (Nchpol/16384) × Nant × (Nant − 1)/(27×26) / (Δt/1 sec)
              ...equation (4)

              Here Nchpol is the sum over all subbands of spectral channels times polarization products:

              Nchpol = Σsb Nchan,i × Npolprod,i

              where Nchan,i is the number of spectral channels in subband i, and Npolprod,i is the number of polarization products for subband i (1 for single polarization [RR or LL], 2 for dual polarization [RR and LL], 4 for full polarization products [RR, RL, LR, LL]). This formula, combined with the maximum data rates given above, imply that observations using the maximum number of channels currently available (16384) will be limited to minimum integration times of ~2 seconds for standard observations, and 0.8 seconds for shared risk observations.

              These data rates are challenging for transfer and analysis. Data may either be downloaded via ftp over the Internet, or mailed on hard drives for large data sets or for those with slow Internet connections (please review the data shipping policy). For users whose science permits, the Archive Access Tool allows some level of frequency averaging in order to decrease data set sizes before ftp; note that the full spectral resolution will be retained in the NRAO archive for all observations.

              Higher time resolutions and data rates are possible in principle but will be considered only through the Resident Shared Risk Observing program.

              Note: The data rate formula given above does not account for the auto-correlations delivered by WIDAR. Precise data rate values can be obtained through the use of the General Observing Setup Tool (GOST), or the Resource Catalog Tool (also known as the Instrument Configuration Tool) of the Observation Preparation Tool (OPT).

              Radio Frequency Interference

              The very wide bandwidths of the VLA mean that radio frequency interference (RFI) will be present in a far larger fraction of current observations than in observations made with the old systems. Considerable effort has gone into making the VLA's electronics as linear as possible, so that the effects of any RFI will remain limited to the actual frequencies at which the RFI exists. Non-linear effects, such as receiver saturation, should occur only for those very unlikely, and usually very brief, times when the emitter is within the antenna primary beam.

              RFI is primarily a problem within the low frequency (C, S, L, and the low-band system) bands, and is most serious in the D configuration. With increasing frequency and increasing resolution comes an increasing fringe rate, which is often very effective in reducing interference to tolerable levels.

              The bands within the tuning range of the VLA which are allocated exclusively to radio astronomy are: 1400–1427 MHz, 1660–1670 MHz, 2690–2700 MHz, 4990–5000 MHz, 10.68–10.7 GHz, 15.35–15.4 GHz, 22.21–22.5 GHz, 23.6–24.0 GHz, 31.3–31.8 GHz, and 42.5–43.5 GHz. No external interference should occur within these bands.

              VLA staff periodically observes the entire radio spectrum with the VLA, from 1.0 through 50.0 GHz with 125 kHz channel resolution, to monitor the ever-changing RFI spectrum. Users concerned about the precise frequencies of strong RFI, and the likelihood of being affected, are encouraged to peruse these plots. To access these plots, or for more information on RFI, including the impact of satellite transmissions, please see the RFI section in the Guide to Observing with the VLA.

              Subarrays

              The continuum subarray option offers two 1 GHz baseband pairs with the 8-bit samplers in up to 3 subarrays, with the same spectral channel and polarization product options as are available for wideband observing. The setup for each subarray is completely independent in terms of observing frequency, polarization products, and integration times.

              When using three subarrays, there are some restrictions on the number of antennas in each subarray. The Baseline Board in the correlator treats each set of 4 antennas independently, using a separate column of correlator chips. With 8 such columns, the correlator can handle up to 8×4 = 32 antennas. The correlator configuration software requires that a given column not be split across subarrays. This does not matter when using only two subarrays, but forces some subtle restrictions when using three. For instance, one cannot observe with 9 antennas in each of 3 subarrays, because 9 antennas requires three columns (two with 4 antennas each, and one with 1 antenna); three subarrays of 9 antennas each would require 3×3 = 9 columns, one more than are actually available. Splitting the array into 10, 9, and 8 antennas is allowed, since the first two subarrays use 3 columns each, while the third uses only two.

              Table 3.8.1 gives four examples of how correlator resources can be split into multiple subarrays. Antennas in each subarray are color-coded: red for subarray 1, green for subarray 2 (if present), and blue for subarray 3 (if present). The last column gives the number of antennas in each subarray (e.g., in the setup shown in the first row, subarray 1 has 10 antennas, subarray 2 has 9 antennas, and subarray 3 has 8 antennas). In all cases a total of 27 antennas are used. The columns are numbered in reverse order (C7 to C0) to match the numbering scheme used on the actual Baseline Boards.

              Table 3.8.1: Some Possible Subarray Options
              Number of antennas correlated using each Baseline Board columnNumber
              of antennas
              C7 C6 C5 C4 C3 C2 C1 C0
              4 4 2 4 4 1 4 4 10 + 9 + 8
              4 4 4 2 4 4 4 1 14 + 13
              4 4 3 4 4 3 3 2 11 + 11 + 5
              4 4 4 4 4 4 3 27

              For more information on subarrays, please see the Subarray section in the Guide to Observing with the VLA.

              Positional Accuracy & Astrometry

              The position of a target can be determined to a small fraction of the synthesized beam, limited by atmospheric phase stability, the proximity of an astrometric calibrator, the calibrator-source cycle time, and the SNR on target.

              In preparation for observing, the a priori position must be known to within the antenna primary beam, except perhaps for mosaicking observations. In the special case of using the phased VLA as a VLBI element, the a priori position must be accurate to within the synthesized beam of the array.

              In post-processing, target positions are typically determined from an image made after phase calibration, i.e., correcting the antenna and atmospheric phases as determined on the reference source. The accuracy of the calibration determines the accuracy of the positions in the image. Note that phase self-calibration imposes the assumed position of the model, i.e., makes the position indeterminate. Therefore, an absolute position cannot be determined after self-calibration, but relative positions between features within a self-calibrated image are valid.

              It may help to think of astrometry as two methods, narrow-field and wide-field.

              Narrow-field astrometry

              In narrow-field astrometry, the target is close to the phase tracking center and the antennas nod every few minutes between the target and a calibrator. Under good conditions of phase stability with accurate antenna positions (baselines), a strong target, a close calibrator with accurately known position, and rapid switching, the accuracy can approach 1–2% of the synthesized beam, with a floor of ~2 milliarcseconds (mas). Even under more typical conditions, 10% of the beam is readily achieved.

              Astrometric calibrators are marked J2000 A in the VLA calibrator list, and have an accuracy of ~2 mas. Other catalogs from the USNO and the VLBA are also useful, but offsets may exist between the VLA and VLBA centroids arising from extended structure in the particular source and the different resolutions of the arrays.

              For studies of proper motion and parallax, the absolute accuracy of a calibrator may be less important than its stability over time. Close, or in-beam calibrators with poor a priori positions, can be used and tied to the ICRF reference frame in the same or separate observations.

              Phase stability can be assessed in real time from the Atmospheric Phase Interferometer (API) at the VLA site, which uses observations of a geostationary satellite at ~12 GHz. Dynamic scheduling uses the API data to run a project under suitable conditions specified by the user. VLBI projects using the phased VLA will typically be fixed date and not dynamically scheduled.

              Wide-field astrometry

              Wide-field astrometry is used to determine the positions of targets within the primary beam, referenced to a calibrator within the beam or close by. In addition to the previous effects, there are distortions as a function of position in the field, from small errors in the Earth Orientation Parameters (EOP) used at correlation time, differential aberration, and phase gradients across the primary beam. With no special effort, the errors build up to roughly one synthesized beam at a separation of ~104 beams from the phase tracking center. Not all these errors are fully understood, and accurate recovery of positions over the full primary beam in the wide-band, wide-field case is a research area. These effects are handled somewhat differently in the post-processing packages. Check with VLA staff for more details via the NRAO Helpdesk.

              Limitations on Imaging Performance

              Imaging Fidelity

              Image fidelity is a measure of the accuracy of the reconstructed sky brightness distribution. A related metric, dynamic range, is a measure of the degree to which imaging artifacts around strong sources are suppressed, which in turn implies a higher fidelity of the on-source reconstruction.

              With conventional external calibration methods, even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is most often the effective phase stability of the telescope due to atmospheric/ionospheric fluctuations, although pointing errors and changes in atmospheric opacity may also be a limiting factor. If a good model of the sky brightness distribution exist (e.g., use of compact structures of sufficient strength—though a good model of resolved sources in the field of view may also be used), standard self-calibration can be counted on to improve the images. At low frequencies, where the dominant phase error is due to ionospheric plasma density fluctuations, more advanced techniques may be required to account for change of ionospheric phase across the field of view. Depending on the underlying nature of the errors, dynamic ranges in the thousands to hundreds of thousands can be achieved using these techniques. With the new WIDAR correlator, with its much greater bandwidths and much higher sensitivities, self-calibration methods can be extended to observations of sources with much lower flux densities than ever possible with the old VLA.

              The choice of image reconstruction algorithm also affects the correctness of the on-source brightness distribution. The CLEAN algorithm is most appropriate for predominantly point-source dominated fields. Extended structure is better reconstructed with multi-resolution and multi-scale algorithms. For high dynamic ranges with wide bandwidths, algorithms that model the sky spectrum as well as the average intensity can yield more accurate reconstructions.

               

              Invisible Structures

              An interferometric array acts as a spatial filter so that, for any given configuration, structures on a scale larger than the fringe spacing of the shortest baseline will be completely absent. Diagnostics of this effect include negative bowls around extended objects, and large-scale stripes in the image. Image reconstruction algorithms such as multi-resolution and multi-scale CLEAN can help to reduce or eliminate these negative bowls, but care must be taken in choosing appropriate scale sizes to work with.

              Table 3.1.1 gives the largest scale visible to each configuration/band combination.

               

              Poorly Sampled Fourier Plane

              Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1 in Documentation.

               

              Sidelobes from non-Deconvolved Sources

              At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth smearing and time-averaging will tend to reduce the effects of these sources, the very best images will require careful imaging of all significant background sources. The deconvolution tasks in AIPS (IMAGR) and CASA (clean) are well suited to this task. Sidelobe confusion is a strong function of observing band—affecting most strongly L and P-band observations—and is rarely a significant problem for observations at frequencies above 4 GHz.

               

              Sidelobes from Strong Sources

              An extension of the previous section is to very strong sources located anywhere in the sky, such as the Sun (especially when a flare is active), or when observing with a few tens of degrees of the very strong sources Cygnus A and Casseopeia A. Image degradation is especially notable at lower frequencies, shorter configurations, and when using narrow-bandwidth observations (especially in spectral-line work) where chromatic aberration cannot be utilized to reduce the disturbances. In general, the only relief is to include the disturbing sources in the imaging, or to observe when these objects are not in the viewable hemisphere.

               

              Wide-band Imaging

              The very wide bandpasses provided by the VLA enable imaging over 2:1 bandwidth ratios: at L, S, and C-bands, the upper frequency is twice that of the lower frequency. It is this wide bandwidth which enables sub-microJy sensitivity.

              In many cases, where the observation goal is a simple detection and there are no strong sources near to the region of interest, standard imaging methods that combine the data from all frequencies into one single image (multi-frequency-synthesis) may suffice. This is because the wideband system makes a much better synthesized beam—especially for longer integrations—than the old, single-frequency beam; considerably reducing the region of sky which is affected by incorrect imaging/deconvolution. A rough rule of thumb is that—provided a strong source is not adjacent to the target zone—if the necessary dynamic range in the image is less than 1000:1 (i.e., the strongest source in the beam is less than 1000 times higher than the noise), a simple wide-band map may suffice.  

              For higher dynamic ranges, complications arise from the fact that the brightness in the field of view dramatically changes as a function of frequency, both due to differing structures in the actual sources in the field of view and due to the attenuation of the sources by the primary beam. One symptom of such problems is the appearance of radial spokes around bright sources, visible above the noise floor, when imaged as described above.  

              The simplest solution is to simply make a number of maps (e.g., one for each subband) which can then be suitably combined after correction for the primary beam shape. But with up to 64 subbands available with the VLA's new correlator, this is not always the optimal approach. Further, images at all bands must be smoothed to the angular resolution at the lowest frequency before any spectral information can be extracted. With a 2:1 bandwidth, the difference in angular resolution across the band will be significant.  

              A better approach is to process all subbands simultaneously, utilizing software which takes into account the possibility of spatially variant spectral index and curvature and knows the instrumentally-imposed attenuation due to the primary beam. Such wideband imaging algorithms are now available within CASA as part of the clean task, and work is underway to integrate them fully with wide-field imaging techniques.

               

              Wide-field Imaging

              Wide-field observing refers primarily to the non-coplanar nature of the VLA when observing in non-snapshot mode. At high angular resolutions and low frequencies, standard imaging methods will produce artifacts around sources away from the phase center.  Faceted imaging (AIPS, CASA) and W-projection (CASA) techniques can be used to solve this problem.

              Another aspect of wide-field observing is the accurate representation of primary beam patterns and their use during imaging. This is relevant only for very high dynamic ranges ( > 10000 ) or when there are very strong, confusing sources at and beyond the half-power point of the primary beam. This problem is worse with a wide-band instrument simply because the size of the primary beam—and the radius at which the half-power point occurs—varies with frequency while there is increased sensitivity out to a wider field of view. Work is underway to commission algorithms that deal with these effects by modeling and correcting for frequency-dependent and rotating primary beams, per antenna, during imaging. Please note, however, that most advanced methods will lead to a significant increase in processing time and may not always be required. Therefore, in the interest of practicality, they should be used only if there is evidence of artifacts without these methods.

              Finally, all of the above effects come into play for mosaicking, another form of wide-field imaging in which data from multiple pointings are combined during or after imaging.

              Calibration and Flux Density Scale

              The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators; and is also available from within the Observation Preparation Tool (OPT).

              Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration—consult the Guide to Observing with the VLA for advice. Over the last several years, accurate source models have been implemented directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 at L, S, C, X, Ku, K, Ka, and Q-bands.

              Since the standard source flux densities are slowly variable, we monitor their flux densities when the array is in its D configuration. As the VLA cannot accurately measure absolute flux densities, the values obtained must be referenced to assumed or calculated standards as described in the next paragraph. Table 3.11.1 shows the flux densities of these sources in January 2012 at the standard VLA bands. The flux density scale for the VLA, from 1 through 50 GHz, is based on emission models of the planet Mars, which is then calibrated to the CMB dipole using WMAP (Wilkinson Microwave Anisotropy Probe) observations (for details see Perley and Butler, 2013a). The source 3C286 (J1331+3030) is known to be non-variable, and has thus been adopted as the prime flux density calibrator source for the VLA. The adopted polynomial expression for the spectral flux density for 3C286 is:

              \[\log(S) = 1.2515 - 0.4605 \log(f) - 0.1715 \log^2(f) + 0.0336 \log^3(f)\]

              where S is the flux density in Jy, and f is the frequency in GHz.

              The absolute accuracy of our flux density scale is estimated to be about 2%. With care, the internal accuracy in flux density bootstrapping is better than 1% at all bands except Q-band, where pointing errors limit bootstrap accuracy to perhaps 3%. Note that such high internal accuracies are only possible in long-duration observations where the antenna gain curves and atmospheric opacity can be directly measured, and where there is good elevation overlap between the target source(s) and the flux density standard calibrator.

              Table 3.11.1: Flux densities (Jy) of Standard Calibrators for January 2012
              Source

              1465 MHz

              2565 MHz4885 MHz8435 MHz14965 MHz22460 MHz36435 MHz43340 MHz
              3C48 = J0137+3309 15.56 9.80 5.39 3.14 1.77 1.19 0.73 0.63
              3C138 = J0521+1638 8.71 6.17 4.02 2.78 1.89 1.46 1.03 0.92
              3C147 = J0542+4951 21.85 13.75 7.59 4.49 2.59 1.77 1.10 0.94
              3C286 = J1331+3030 14.90 10.03 7.34 5.09 3.39 2.52 1.75 1.53
              3C295 = J1411+5212 22.15 12.95 6.41 3.34 1.62 0.957 0.507 0.403
              NGC7027 1.62 3.59 5.38 5.79 5.62 5.42 5.18 5.04

              The sources 3C48, 3C147, and 3C138 are all slowly variable. VLA staff monitor these variations on timescales of a year or two;  suitable polynomial coefficients are determined for them which should allow accurate flux density bootstrapping. These coefficients are updated approximately every other year and are used in the AIPS task SETJY and in the CASA task setjy.

              The VLA antennas have elevation-dependent gain variations which are important to account for at the four highest frequency bands. Gain curves are determined by VLA staff, and the necessary corrections are applied to the visibility data when these data are downloaded from the archive. Additionally, atmospheric opacity will also cause an elevation-dependent gain which is also particularly notable at these four highest frequency bands. Currently, we do not have an atmospheric opacity monitoring procedure; users should utilize the appropriate tasks available in both AIPS and CASA to estimate and correct for the opacity using ground-based weather data. Correction of these gain dependencies, plus regular calibration using a nearby phase calibrator, should enable good amplitude gain calibration for most users. Please note that extraordinary attenuation by clouds can only be approximately corrected for by regular observation of a nearby calibrator.

              A better procedure for removing elevation gain dependencies uses the AIPS task ELINT. This task will generate a 2nd order polynomial gain correction utilizing your own calibrator observations. This will remove both the antenna and opacity gain variations, and has the decided advantage of not utilizing opacity models or possibly outdated antenna gain curves. Use of this procedure is only practical if your observations span a wide range in elevation.

              By far the most important gain variation effect is that due to pointing. Daytime observations on sunny days can suffer pointing errors, primarily in elevation, of up to one arcminute. This effect can be largely removed by utilizing the referenced pointing procedure which determines the pointing offset of a nearby calibrator. This offset is then applied to subsequent target source observations. It is recommended that this local offset be determined at least hourly, utilizing an object within 15 degrees of the target source—preferentially at an earlier R.A.. Studies show that the maximum pointing error will be reduced to about 7 arcseconds or better. VLA staff continue to work on improving this essential methodology.

              The VLA's post-amplifiers are not temperature stabilized and exhibit significant gain changes between night and day, particularly at the four highest frequency bands. Changes as large as 30% have been seen between night and day in calm, clear conditions. These gain changes, and others caused by possible changes in attenuator settings, are monitored and will be removed with excellent accuracy by application of the internal calibration signal, whose results are recorded in the switched power table (SY table in AIPS). These corrections are not applied by default—users who wish to correct for these gain changes must utilize the appropriate tasks in AIPS or CASA. For the most accurate flux density bootstrapping, this table must be applied to the visibility data before calibration. Gain bootstrapping better than 1% can be accomplished for the 8-bit sampler system after application of the switched power data. For the 3-bit system there is an additional complication, as the values of the switched power data are sensitive to the total power as well as the system gain. VLA staff are currently working on a methodology to remove the total power dependency. Not applying the switched power data will reduce bootstrapping accuracy to perhaps 10%, and possibly worse, if the observation of the flux density calibrator is not close in time to the local complex gain calibrator (amplitude and phase).

              Complex Gain Calibration

              General Guidelines for Gain Calibration

              Adequate gain calibration is a complicated function of source-calibrator separation, frequency, array scale, and weather. Since what defines adequate for some experiments is completely inadequate for others, it is difficult to define simple guidelines to ensure adequate phase calibration. However, some general statements remain valid most of the time. These are given below.

              • Under decent conditions with no thunderstorms or ionospheric storms, tropospheric effects dominate at frequencies higher than about 4 GHz; ionospheric effects dominate at frequencies lower than about 4 GHz.
              • Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the C and D configurations at L and S-bands, and in the D configuration at X and C-bands. For these cases, calibration need only be done to track instrumental changes—a couple of times per hour is usually sufficient.
              • If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies or 15 minutes at high frequencies in order to track pointing or other effects that might influence the amplitude scale. The enhanced sensitivity of the VLA guarantees, for full-band continuum observations, that every field will have enough background sources to enable phase self-calibration at L and S-bands. At higher frequencies, the background sky is not sufficient, and only the flux of the target source itself will be available.
              • The smaller the source-calibrator angular separation, the better. In deciding between a nearby calibrator with an S code in the calibrator database, and a more distant calibrator with a P code, the nearby calibrator is usually the better choice. A detailed description of calibrator codes is available in the calibrator list.
              • In clear and calm conditions, most notably in the summer, phase stability often deteriorates dramatically after about 10AM due to small-scale convective cells set up by solar heating. Observers should consider a more rapid calibration cycle for observations between this time and a couple hours after sundown.
              • At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) (below).

               

              Rapid Phase Calibration and the Atmospheric Phase Interferometer

              For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source-calibrator observing cycles as short as 40 seconds can be used for very small source-calibrator separations. Observing efficiency declines, however, for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the old VLA this was known as fast-switching; for the upgraded VLA, it is a loop of source-calibrator scans with short scan length. This technique stops tropospheric phase variations at an effective baseline length of ∼vat/2, where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec) and t is the total switching time. Short source-calibration scans have been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Note that at a typical wind velocity in the compact D-configuration, this effective baseline length is the same as—or larger than—the longest baseline in the array and it is not worth the increased overhead of short cycle times. Under these circumstances, it is sufficient to calibrate every 5-10 minutes to track the instrumental changes. The fast switching technique will not work in bad weather (such as rain showers or when there are well-developed convection cells (thunderstorms)). It is also important to correctly specify the required tropospheric phase stability as measured by the Atmospheric Phase Interferometer at observe time (see below).

              Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from References 9 and 10 in Documentation.  See also the High Frequency Observing guide for additional recommendations on observing at high frequencies.

              An Atmospheric Phase Interferometer (API) is used to continuously measure the tropospheric contribution to the interferometric phase. The API uses an interferometer of two, 1.5 meter antennas, separated by 300 meters, observing an 11.7 GHz beacon from a geostationary satellite. The API data are heavily used for the dynamic scheduling of the VLA.

              Characteristic seasonal averages are represented in Table 3.12.1 below:

              Table 3.12.1: Seasonal API/wind values at the VLA
              Month

              API (night)
              [deg]

              API (median)
              [deg]

              API (day)
              [deg]

              Wind (night)
              [m/s]

              Wind (median)
              [m/s]

              Wind (day)
              [m/s]

              January 2.3 2.8 3.6 1.6 1.9 2.3
              February 2.9 3.4 4.5 4.0 4.3 4.5
              March 2.8 3.7 5.5 3.4 3.9 4.7
              April 3.3 4.5 6.2 5.3 5.5 5.8
              May 2.9 4.6 6.7 2.6 3.2 3.7
              June 3.8 5.5 7.4 2.5 3.9 6.3
              July 6.2 8.3 10.5 2.9 2.9 3.0
              August 5.4 7.1 11.3 1.7 2.3 3.0
              September 5.2 6.6 8.8 2.3 3.0 3.6
              October 4.2 5.3 7.4 2.3 2.9 3.7
              November 2.6 3.0 4.0 1.2 2.5 1.6
              December 2.8 3.2 4.1 1.2 1.6 2.7

              Day indicates sunrise to sunset values; night indicates sunset to sunrise values.

                Polarization

                For projects requiring imaging in Stokes Q and U, the instrumental polarization should be determined through observations of a bright calibrator source spread over a range in parallactic angle. The complex gain calibrator chosen for the observations can also double as a polarization calibrator, provided it is at a declination where it moves through enough parallactic angle during the observation (roughly Dec 15–50 degrees during a several hour track). The minimum condition that will enable accurate polarization calibration from a polarized source, in particular with unknown polarization, is three observations of a bright source spanning at least 60 degrees in parallactic angle (schedule four scans, if possible, in case one is lost); if at all possible, it is strongly recommended that five or more observations covering 100 degrees (or more) of parallactic angle in roughly uniform steps be run. If a bright, unpolarized, unresolved source is available, and known to have very low polarization, then a single scan will suffice to determine the leakage terms. The accuracy of polarization calibration is generally better than 0.5% for small objects as compared to the antenna beam size. At least one observation of 3C286 or 3C138 is required to fix the absolute position angle of polarized emission; 3C48 can be used at frequencies of ~3 GHz and higher, or 3C147 at frequencies over ~10 GHz. Note that 3C138 is variable—the polarization properties are known to be changing significantly over time, most notably at the higher frequencies (for details see Perley and Butler (2013b)).

                More information on polarization calibration strategy can be found in the Polarimetry section in the Guide to Observing with the VLA.

                 

                VLBI Observations

                The VLA can participate in VLBI observations with the VLBA. Presently, this is only allowed in phased array mode (single dish is only available through the VLBA Resident Shared Risk Observing program) with restricted WIDAR correlator configurations. Also note that, currently, P-band cannot be phased. For more details see the VLBI at the VLA documentation. In phased array mode the program TelCal derives the antenna-based delay and phase corrections needed for antenna phasing in real time. This correction is applied to the antenna signals before they are summed, requantized to 2-bits, and recorded in VDIF format on the Mark5C disk at the VLA site. The disk(s) are then transported to Socorro, NM and correlated on the DiFX correlator with other VLBI stations which participated in the observation. Standard VLA data, i.e., correlations between VLA antennas, are also archived in the NRAO science data archive.

                Snapshots

                The two-dimensional geometry of the VLA allows a snapshot mode whereby short observations can be used to image relatively bright, unconfused sources. This mode is ideal for survey work where the sensitivity requirements are modest.

                Single snapshots with good phase stability of strong sources should give dynamic ranges of a few hundred. Note that because the snapshot synthesized beam contains high sidelobes, the effects of background confusing sources are much worse than for full syntheses, especially at 20 cm and longer wavelengths in the D configuration. For instance, at 20 cm, a single snapshot will give a limiting noise of about 0.2 mJy. This level can be reduced by taking multiple snapshots separated by at least one hour. The deconvolution of the data is necessary to remove the effects of background sources. Before considering snapshot observations at 20 cm, users should first determine if the goals desired can be achieved with the existing Faint Images of the Radio Sky at Twenty-centimeters survey (FIRST, http://sundog.stsci.edu/top.html) (B configuration) or the NRAO VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) (D configuration, all-sky).

                Shadowing and Cross-Talk

                Observations at low elevation in the C and D configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling when using the default inputs with CASA tasks importasdm and importevla. AIPS task UVFLG can be used to flag VLA data based on shadowing, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other subarrays. The CASA task flagdata can also be used to flag data based on shadowing. For more information on shadowing, please see the Antenna Shadowing section in the Guide to Observing with the VLA.

                Cross-talk is an effect in which signals from one antenna are picked up by an adjacent antenna, causing an erroneous correlation. This effect is important at low frequencies in compact configurations. Careful examination of the visibilities is necessary to identify and remove this form of interference. The affected data would show time-variable high-amplitude points.

                Combining Configurations and Mosaicking

                Any single VLA configuration will allow accurate imaging of a range of spatial scales determined by the shortest and longest baselines. For extended and structured objects, it may be required to obtain observations in multiple array configurations. It is advisable that the frequencies used be the same for all configurations to be combined. The ideal combination of arrays results in a uv-plane with all cells equally filled by uv-points. To first order, this can be achieved by using the beam sizes of the individual arrays to inversely scale the on-source integration time. This approach is equivalent to achieving the same surface brightness sensitivity for all arrays on all scales. For the VLA, observations in the different configurations generate beam sizes that decrease by factors of three, i.e., C configuration generates a three times smaller beam than D configuration, B three times smaller than C, and A three times smaller than B. Thus, on-source integrations would increase by about an order of magnitude between each array. Such a drastic increase is very expensive and, in fact, not necessary since some spatial scales are common to more than a single array, which is equivalent to some uv-cells being filled more than others. The best way to fill the uv-plane depends on many factors such as declination of the source, LST time of the observation, and bandwidth.

                Experience shows for the VLA that a factor of about three in on-source integration time for the different array configurations works well for most experiments. For example, a 20min on-source time in D, 1hr in C, 3hrs in B, and 9hrs in A should produce a decent map. Using large bandwidths and multi-frequency synthesis will broaden all uv tracks radially and one may need even less array configurations or shorter integration times between the different arrays.

                Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicking. The VLA has no limit on the number of pointings for each mosaic. Typically hexagonal, rectangular, or individual pointing patterns are used and the overlap regions will result in an improved rms over each individual pointing. Given the many, potentially short observations, it is important to obey the data rate limits outlined in the Time Resolution and Data Rates Section. On-the-fly (OTF) mosaics, i.e. dumping the data while moving the telescopes across the source, is also available.

                Time-variable structures, such as the nuclei of radio galaxies and quasars, cause special, but manageable, problems. See the article by Mark Holdaway in Reference 2 of the Documentation for more information.

                Guidelines for mosaicking with the VLA are given in the Guide to Observing with the VLA.

                Pulsar Observing

                The VLA can be used for several kinds of pulsar observing: phase-binning using the WIDAR correlator, using the phased-array for single-beam pulsar processing in either search or fold modes, or simply standard imaging mode with fast integrations. Any of these types of pulsar observing are considered Resident Shared Risk Observing (RSRO), and participants in this program are expected to work closely with NRAO staff. Of these, phased-array and fast-dump observing are significantly more mature than phase binning. For any questions not addressed here regarding the capabilities of these observing modes, please contact the NRAO Helpdesk.

                Phased-array pulsar processing

                The "Y" Ultimate Pulsar Processing Instrument (YUPPI) is a suite of software that runs in the correlator backend (CBE) computer cluster and can process a single-beam phased/summed-array data stream for pulsar observations in real time, into either folded profiles or search mode (filterbank) output. Coherent dedispersion can be optionally applied in either mode.

                In this mode, the voltage data streams from each antenna are divided into a number of frequency subbands within the correlator, then summed and requantized before being output to the cluster for pulsar processing. The limitation on bandwidth comes primarily from the available network connections between the correlator and cluster. In all cases, a maximum of 64 subbands total can be processed. Depending on the number of bits chosen, this results in the following total bandwidth constraints:

                Table 3.18.1: Pulsar Observing Bandwidth Constraints

                Subband bandwidth

                Subband quantization

                Max total bandwidth

                Samplers

                32 MHz 8 bits 2048 MHz 8-bit
                64 MHz 4 bits 4096 MHz 3-bit
                128 MHz 2 bits 8192 MHz 3-bit
                <32 MHz 8 bits 64*BWsub 8-bit

                As described in the VLA Frequency Bands and Tunability section, the 8-bit samplers provide two independently tunable 1 GHz IFs, while the 3-bit samplers provide four tunable 2 GHz IFs. 

                The pulsar-specific processing is done in real time using the DSPSR software package and, in principle, any processing option supported by DSPSR can be used, although this will be constrained by the real-time computing power available in the cluster. In general, each subband can be divided into an arbitrary (2n) number of channels; 1 (summed), 2 or 4 detected polarization products can be output; and coherent dedispersion can be enabled or not.

                Fold mode

                In fold mode, the data are averaged modulo a known pulsar ephemeris (provided via a standard TEMPO/TEMPO2 "par file") into pulse profiles. The data can also be folded at a constant topocentric period, for example at 10 Hz to detect the injected noise cal signal. Fold integration times as short as 1 second have been tested. Up to 16384 profile bins can be used. The data are recorded in PSRFITS format using the standard 16-bit data encoding. This means the final output data rate is given by:

                Data rate = 2 bytes × Nsubband × Nchannel × Nbin × Npoln / Tint

                If the desired data rate exceeds ~25 MB/s, additional testing ahead of time may be required.

                Search mode

                In search mode, the data are simply detected and averaged over a specified amount of time before being output to disk, resulting in a filterbank data array (power vs time and frequency). Coherent dedispersion at a known DM can optionally be enabled for this.  Data can be recorded using 2, 4, 8, 16 or 32 bits, resulting in a final data rate of:

                Data rate = (Nbit/8) bytes × Nsubband × Nchannel × Npoln / Tint

                The maximum sustained output rate in this mode should be kept less than ~400 MB/s.

                Gated or binned visibilities

                The WIDAR correlator has the capability to internally integrate (fold) visibilities into 1 or more pulse phase bins. Constraints and trade-offs on number of bins, bin width, pulse period, bandwidth, and integration time are currently not well quantified. This mode is likely to require significant development work to become usable.

                Fast-dump visibilities

                While not specifically a pulsar mode, standard visibility data can be dumped as fast as 5 ms, which may be sufficient for imaging of slow pulsars.  See the Time Resolution and Data Rates section for more details.

                WIDAR Correlator

                Introduction

                The correlator configurations offered for general observing may be divided into three basic modes: wideband, spectral line, and subarrays. The possible setups are also subject to the integration time and data rate restrictions outlined in the section on Time Resolution and Data Rates. The possibilities and restrictions are embodied in the General Observing Setup Tool (GOST) and in the Resources section of the Proposal Submission Tool (PST), which must be used to define the correlator configuration for General Observing (GO) and Shared Risk Observing (SRO) proposals.

                Note that phased array configurations are only allowed as part of VLBI experiments (see the section on VLBI Observations) or as Resident Shared Risk observations.

                Wideband and spectral line observing modes with the WIDAR correlator are described below. For the subarray mode, we refer to the Subarrays section of the OSS.

                 

                WIDAR Correlator: Wideband Observing

                The wideband observing setups provide the widest possible bandwidth for a given observing band, with channel spacing depending on the number of polarization products as listed in the following table 4.1.1:

                 

                Table 4.1.1: Wideband & Subarray Correlator Options
                (all but P and L-bands)
                Polarization productsChannel spacing
                Full (RR, RL, LR, LL) 2 MHz
                Dual (RR and LL) 1 MHz
                Single (RR or LL) 0.5 MHz

                 

                8-bit wideband setups are available for all observing bands, providing a total of 2 GHz of bandwidth per polarization (1 GHz per polarization at L-band, and 256 MHz per polarization at P-band). 3-bit setups are available for all bands above S-band, providing total bandwidths per polarization of 4 GHz (C/X-bands), 6 GHz (Ku-band), or 8 GHz (K/Ka/Q-bands). In all cases, except for P and L-band, each of the subbands is 128 MHz wide. At L-band the default is 64 MHz/subband, yielding channels twice as narrow as those listed in the table above, while at P-band the default is 16 MHz/subband, resulting in 125 kHz channel spacing.

                In many frequency bands, the total processed bandwidth is less than that delivered by the front-end. In those cases, the observer may independently tune two 1 GHz baseband pairs when using the 8-bit samplers, or four 2 GHz baseband pairs when using the 3-bit samplers, or choose to have a mix 8-bit and 3-bit samplers. The tuning restrictions are described in the section on VLA Frequency Bands and Tunability, and the 8-bit and 3-bit samplers are described in the section on VLA Samplers.

                 

                WIDAR Correlator: Spectral Line Observing



                Basebands and Subbands

                Observers have access to very flexible correlator configurations using up to 64 subbands in up to 4 basebands sampled with the 8-bit and/or the 3-bit samplers. These capabilities may be summarized as follows:

                • Two 1 GHz baseband pairs using the 8-bit samplers, or four 2 GHz baseband pairs using the 3-bit samplers, independently tunable within the limits outlined in the section on VLA Frequency Bands and Tunability. The 8-bit baseband pairs are referred to as A0/C0 and B0/D0, while the 3-bit samplers are A1/C1, A2/C2, B1/D1, and B2/D2. The AC/BD nomenclature corresponds to that of the IF pairs in the pre-expansion VLA.
                • Up to 16 subband pairs (spectral windows) in each 3-bit baseband pair, and up to 32 subbands in each 8-bit baseband pair, for a total of up to 64 subbands in any correlator configuration:
                  • Tuning, bandwidth, number of polarization products, and number of channels can be selected independently for each subband;
                  • All subbands must share the same integration time;
                  • No part of a subband can cross a 128 MHz boundary;
                  • Subband bandwidths can be 128, 64, 32, …, 0.03125 MHz (128 / 2n, n=0, 1, …, 12).
                • The sum over subbands of channels times polarization products is limited to 16384 (without recirculation):
                  • These may be spread flexibly over subbands and polarization products, in multiples of 64: 64, 128, 192, 256, 384, …, 16384 cross-correlation products;
                  • Recirculation may be used to increase the number of channels per subband for subbands narrower than 128 MHz. Baseline Board stacking may be used to increase the number of channels per subband for setups requiring less than 64 subbands;
                  • Assigning many channels to a given subband may reduce the total bandwidth and/or the total number of subbands available.

                The remainder of this section discusses the various limitations in more detail, including some examples to show how they come up in practice.

                 

                Subband Tuning Restrictions

                Each subband may be placed anywhere within a baseband, with the caveat that no subband may cross a 128 MHz boundary. Mathematically:

                νBB0 + n×128 MHz <= νsbLow <= νsbHigh <= νBB0 + (n+1)×128 MHz

                where:

                νBB0 the lower frequency edge of the baseband;
                n= 0, 1, …, 7 (, …, 15) (any integer between 0 and 7 for 8-bit, between 0 and 15 for 3-bit);
                νsbLow                                   
                the lower edge of the subband
                (the subband center frequency minus half the subband bandwidth);
                νsbHigh                                 
                the upper edge of the subband
                (the subband center frequency plus half the subband bandwidth).

                For example, if the baseband were tuned to cover 10000–11024 MHz, one could place a 64 MHz subband to cover 10570–10634 MHz, but not to cover 10600–10664 MHz because that would cross the 128 MHz boundary at 10640 MHz. Note in particular that the center of a baseband is a boundary and no line should be observed at the baseband center.

                The figure below illustrates these restrictions:

                Correlator configuration figure: bandpass8jul12.png

                The black curve shows the analog filter response for an 8-bit baseband covering 1024 MHz, starting at νBB0. The dashed blue vertical lines show the 128 MHz boundaries; no subband can cross those boundaries and 128 MHz subbands are thus constrained to cover a region between two of those boundaries, with no finer tuning being possible. Narrower subbands, like the 64 MHz subband shown here in red, can be shifted around arbitrarily within one of the 128 MHz slots, but cannot cross any of these boundaries. (The dotted vertical red lines show the boundaries of the 64 MHz subband, while the solid curve shows an illustrative line within the subband.)

                The analog filter shape defining the baseband rolls off severely at one edge of the baseband, so the 128 MHz slot at that edge has reduced sensitivity. The baseband edge is at the lowest sky frequency in the baseband when using the upper sideband, and at the highest sky frequency in the baseband when using the lower sideband.

                 

                Subband Bandwidths and the Digital Filter Response

                The bandwidth for each subband may be selected independently, and can be any of 128/2n MHz, for n= 0, 1, …, 12: 128, 64, 32, 16, 8, 4, 2, or 1 MHz, or 500, 250, 125, 62.5, or 31.25 kHz.

                The usable portion of the subband is set by three effects. First, as discussed above, the analog filters which define the baseband are not perfect, leading to lower sensitivity in the 128 MHz near the baseband edge for the 8-bit samplers.

                Second, because the digital filters are not infinitely sharp, the rejected sideband leaks in at both edges of the subband. This leads to additional (aliased) noise, with a factor ~2 increase in the noise at the subband edges, dropping to a few percent within a few percent of the subband edge. The precise filter shape and noise increase is a complex but predictable function of the subband bandwidth (sbBW) and the subband tuning.

                The third effect stems from the offset frequencies used for sideband rejection in the WIDAR correlator. The local oscillators at the individual antennas are tuned to slightly different frequencies, with those offsets taken out in the correlator. This means that each antenna observes a slightly different sky frequency, and thus some baselines will not give an interesting correlation near one edge of the subband. The maximum frequency shift is currently set to 32×f0, with the fundamental f0 being set to f0 = max(25.6 kHz×sbBW/128 MHz, 100 Hz). Here sbBW is the smallest subband bandwidth within the baseband. For the wider subband bandwidths the maximum frequency shift corresponds to <1% of that bandwidth, but for narrower subbands the effect can be severe. For instance, a 31.25 kHz subband has f0 = 100 Hz, and a maximum frequency shift of 3.2 kHz—10% of the subband may be lost on some baselines.

                 

                Spectral Channels and Polarization Products

                Each subband, without recirculation enabled, can have a different number of channels and polarization products, subject to two limitations:

                1. For the ithsubband, the number of spectral channels can be:
                  • 64 nBlBP,i with full polarization products (RR,RL,LR,LL)
                  • 128 nBlBP,i with dual polarization products (RR and LL)
                  • 256 nBlBP,i with a single polarization product (RR or LL)
                  Here nBlBP,i= 1, 2, 3, 4, 5, …, 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.
                2. The sum over all subbands of nBlBP,i must be less than or equal to 64, the number of Baseline Board pairs in the correlator. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64 × 256 = 16,384 (without recirculation).

                Baseline Boards are the boards in the WIDAR correlator where the actual cross-multiplications are done. There are 128 Baseline Boards arranged as 64 Baseline Board pairs (BlBPs). The limitations given here correspond to the capabilities of the individual boards and the finite number of boards the correlator has.

                Limitation #1 corresponds to table 4.2.1 of the options for subband bandwidth and spectral resolution when using nBlBP Baseline Board pairs for a subband:

                Table 4.2.1: Subband Bandwidth and Spectral Resolution Options (without recirculation)
                Subband bandwidth &
                total velocity coverage
                Full polarization products
                (RR, RL, LR, LL)
                64nBlBP spectral channels

                Channel spacing:
                Dual polarization products
                (RR and LL)
                128nBlBP spectral channels
                Channel spacing:
                Single polarization product
                (RR or LL)
                256nBlBP spectral channels

                Channel spacing:
                128 MHz 38400/νGHz km/s 2000/nBlBP kHz 600/nBlBPGHz km/s 1000/nBlBP kHz 300/nBlBPGHz km/s 500/nBlBP kHz 150/nBlBPGHz km/s
                64 19200 1000 / nBlBP 300 / nBlBP 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP
                32 9600 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP
                16 4800 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP
                8 2400 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP
                4 1200 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 /n BlBP
                2 600 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP
                1 300 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP
                0.5 150 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP
                0.25 75 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP
                0.125 37.5 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP
                0.0625 18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
                0.0325 9.375 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP 0.122 / nBlBP 0.037 / nBlBP

                Subband bandwidth and spectral resolution options. Note that the table entries refer to the spacing between spectral channels—that spacing is before any frequency smoothing, so these channels are not independent.

                • nBlBP is the number of Baseline Board Pairs assigned to the subband.
                • Each subband may have a different number of spectral channels and polarization products, and each may be tuned independently.
                • There can be at most 16 subbands per baseband, and nBlBP must be an integer: 1, 2, 3, 4, 5, …, 64.
                • The sum of nBlBP over all subbands must be less than or equal to 64.
                • Use of more than one BlBP for a subband may further restrict the number of subbands available in one or more of the basebands; see text for details.

                Here are four examples of allowed general observing setups which use all 64 BlBPs to produce the maximum number of channels times polarization products:

                Table 4.2.2: Example BlBP Setups
                Baseband Subband

                Pol'n
                Products

                Spectral
                channels

                nBlBP
                Example 1 A0/C0 sb0 RR 16384 64
                Example 2 A0/C0 sb0 RR 8192 32
                A0/C0 sb1 RR, LL 1024 8
                A0/C0 sb2 RR, LL 512 4
                B0/D0 sb0 RR, LL 2048 16
                B0/D0 sb1 RR,RL,LR,LL 256 4
                Example 3 A0/C0 sb0 RR 8192 32
                A0/C0 sb1 LL 1024 4
                A0/C0 sb2 RR, LL 1024 8
                A0/C0 sb3 RR,RL,LR,LL 1024 16
                A0/C0 sb4 RR,RL,LR,LL 256 4
                Example 4 A0/C0 sb0-5 RR,RL,LR,LL 64 6 x 1
                A0/C0 sb6 RR, LL 3840 1 x 30
                A0/C0 sb7 RR 768 1 x 3
                A0/C0 sb8 RR,RL,LR,LL 192 1 x 3
                B0/D0 sb0-2 RR,RL,LR,LL 64 3 x 1
                B0/D0 sb3 LL 768 1 x 3
                B0/D0 sb4 RR, LL 2048 1 x 16

                 

                Recirculation

                Recirculation is a term to describe the method of increasing the number of spectral channels in a subband using correlator software, as opposed to Baseline Board stacking, which uses correlator hardware (see below). Recirculation is presently achieved by limiting the subband bandwidth and is available only for subbands less than 128 MHz wide. When limiting the bandwidth in a subband, the correlator software can be directed to use the remaining CPU cycles on a Baseline Board pair to obtain more lags in factors of two, running the data through the board for a second, third, etc., time; hence recirculation.

                At some time in the future an alternative method of recirculation, using the extra CPU cycles freed up by increasing the integration time, will be made available. Recirculation, by limiting of the subband bandwidth to increase the number of channels in factors of two, was used in the pre-expansion VLA correlator.

                 

                Recirculation vs. Baseline Board Stacking

                When faced with the choice between recirculation and Baseline Board stacking to increase the number of channels in a subband, we recommend recirculation for subbands narrower than 128 MHz and is supported in observatory software (GOST, OPT). For subbands of 128 MHz, Baseline Board stacking should be utilized to increase the number of channels.

                The present implementation of recirculation is that, for each halving of the subband bandwidth, the number of channels in the subband may be doubled without having to trade off the use of other subbands. Because recirculation is achieved by limiting the subband bandwidth, it is not supported for 128 MHz subbands; therefore for 64 MHz subbands a factor 2 recirculation is supported, etc. The maximum recirculation factor for a subband is 128/(subband bandwidth in MHz) and, of course, subject to other configuration restrictions such as data rate.

                The juggling between the requested number of channels, subband bandwidth, and the available number of Baseline Board pairs is dependent on the science goals and not easily formulated in a standard answer. However, if subbands of less than 128 MHz are used, recirculation becomes an option for setups that can also be achieved with Baseline Board stacking. In such cases, we suggest to use recirculation where possible and within the General Observing or Shared Risk requirements. This frees up unused Baseline Board pairs for other use; alternatively, one becomes less dependent on all Baseline Board pairs being in working order.

                Recirculation with factors 8 to 64 is designated Shared Risk, and recirculation with factors over 64 as Resident Shared Risk. The latter choice may have severe implications for the sensitivity as visibility integration time is used as trade-off. Ask the NRAO Helpdesk for more details.

                 

                Baseline Board Stacking

                As opposed to recirculation, which increases the number of channels in a subband by exploiting otherwise unused CPU resources, Baseline Board stacking adds more channels to a subband by adding correlator hardware resources, i.e., using up more Baseline Board pairs. Using Baseline Board stacking may therefore limit the number of subbands available in one or more of the basebands. Understanding how this works requires understanding some of the details of the correlator hardware. That understanding is built into the GOST, and observers may simply use that tool to find out whether their particular setup will, in fact, work. But the results can be confusing without some understanding of the hardware constraints from which they arise. These hardware constraints are complex, and most observers will not need to understand these details. The following section is for those who are attempting complex line experiments and who find the GOST or the RCT restricting the number of subbands and/or channels they can use in unexpected ways. Most observers can skip the following section.

                Baseline Board Stacking and Correlator Use

                First let us consider how the correlator hardware is organized. The cross-multiplications in the WIDAR correlator are spread across 64 Baseline Board pairs (BlBP), arranged into 4 quadrants of 16 BlBP each. Each baseband is connected directly to one of those quadrants. In the simplest mode, each of the 16 BlBP of a quadrant handles the correlations for one of the 16 subbands of the corresponding baseband. Four basebands and four quadrants are required to handle the full 8 GHz of bandwidth per polarization provided by the 3-bit (wideband) samplers: 8 GHz is split into four basebands of 2 GHz each, with each baseband fed into a different BlBP quadrant. Each BlBP in that quadrant handles a subband of maximum bandwidth 128 MHz, so 16 BlBP handles 16 subbands for a total of 16×128 MHz = 2048 MHz.

                A single BlBP produces 256 cross-correlations per baseline for a single subband, which can be used for a single polarization product (e.g., RR or LL with 256 spectral channels), or two (RR and LL with 128 spectral channels each), or four (RR, RL, LR, and LL with 64 spectral channels each).

                 

                When using the 8-bit samplers, the total bandwidth is only 2 GHz per polarization, split into two basebands of 1 GHz each. The simplest continuum setup uses only two quadrants, since there are only two basebands; and only 8 subbands are required to span the 8×128 MHz = 1024 MHz of each baseband. Three-quarters of the correlator BlBP hardware remain unused.

                 

                The spectral line mode allows access to these extra correlator resources through Baseline Board stacking: using multiple BlBPs to process the same subband and produce more cross-correlations for that subband. This is done using crossbar switches which make the data for a single subband available to several BlBPs. Those BlBPs can then be used to produce more spectral channels for that subband, with n BlBPs producing 256×n cross-correlations per baseline. The limit on the total number of cross-correlations (16384) stems from the total number of BlBPs (64): 64×256 = 16384.

                Unfortunately, completely flexible crossbar switches are expensive and could not be implemented in the VLA's correlator. This means that one cannot route a given subband to a randomly-chosen BlBP. The routings which are possible, are as follows:

                1. A subband in a baseband can be routed to any BlBP within the corresponding quadrant.
                2. Data coming into a given BlBP in one quadrant, can be routed to the corresponding BlBP in any other quadrant.

                Routing option #1 means that one could use all the BlBPs within a quadrant to correlate a single subband, yielding 16×256 = 4096 cross-correlations for that subband:

                 

                Routing option #2 means that one could use the BlBPs in all 4 quadrants to correlate a single subband. One simple case would use 4 BlBPs to correlate each of the 16 subbands in a single baseband, yielding 4×256 = 1024 cross-correlations for each of those subbands. Note that in this case, no BlBPs are left to correlate any data from the second baseband.

                Using routing option #2 does come with a subtle cost: assigning a BlBP in quadrant X to correlate a subband corresponding to quadrant Y removes that BlBP from use in the baseband corresponding to quadrant X…and therefore also removes the corresponding subband in that baseband. So, getting more channels for a subband in one baseband may prevent the use of a subband in a different baseband. To take a simple example, consider an experiment where one wishes to observe a single line in dual polarization with 512 channels (requiring 4 BlBPs), plus as much continuum bandwidth as possible. Naively, one would say there are 16 subbands in each baseband; one is used for the spectral line, so that leaves 16+15 = 31 subbands, and with the widest subband bandwidth (128 MHz) the total available continuum should be 31×128 MHz = 3968 MHz per polarization. Actually, however, there are only 15+15 subbands available, or 30×128 MHz = 3840 MHz per polarization, because the spectral line subband has eaten one BlBP corresponding to the other baseband:

                 If the same spectral line required twice as many channels, this will result in the loss of two subbands in both of the basebands:

                In some cases one may want to use a different routing to use up subbands in one baseband in preference to another. For instance, the same spectral line setup (2048 cross-correlations for a single spectral line subband, plus as much continuum as possible) could be set up to allow 13 continuum subbands in the A0/C0 baseband, and the full 16 continuum subbands in B0/D0:

                Understanding these confusing constraints can help observers set up the VLA more effectively to achieve their scientific goals. For instance, in a mixed line+continuum experiment, it works best to use the resource tools to set up the baseband tunings and subband channelization for the most important lines first, then add the desired continuum, and then see what correlator resources remain for any lines of secondary interest.

                The above examples all use BlBP pair stacking in powers of 2, but this is not required. To give some idea of more complex possibilities, the following tables (4.2.3 and 4.2.4) give two examples of other possible configurations. The RCT display shows how the Baseline Boards are used to process the individual subbands. The cyan boxes (shaded when printed out in black and white) show the Baseline Boards used to process data from baseband A0/C0, while the yellow boxes show Baseline Boards used to process data from baseband B0/D0.

                Table 4.2.3: Complex Configuration Example #1
                BasebandSubbandPol'n productsSpectral channelsnBlBP
                A0/C0 sb0 RR 10240 40
                A0/C0 sb1 LL 768 3
                A0/C0 sb2 RR,LL 2176 17
                B0/D0 sb0 RR 256 1
                B0/D0 sb1 RR,LL 384 3
                RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

                 

                Table 4.2.4: Complex Configuration Example #2
                BasebandSubbandPol'n productsSpectral channelsnBlBP
                A0/C0 sb0 RR 4352 17
                A0/C0 sb1 RR, LL 1152 9
                B0/D0 sb0 RR,RL,LR,LL 192 3
                B0/D0 sb1 RR, LL 4480 35
                RCT display: corr-cfg-fig:sro2_8bit_ac17+9_bd3+35

                 

                The individual subbands can have different bandwidths, and those bandwidths may be chosen completely independently of the number of spectral channels in each subband. So, for instance, a subband with a bandwidth of 2 MHz and 1152 spectral channels would have a channel separation of 2 MHz/1152 = 1.736 kHz; but the observer could equally well choose a bandwidth of 64 MHz for that subband, leading to a channel separation of 64 MHz/1152 = 55.56 kHz.

                Use of the 3-bit samplers further extends the possibilities.  Here is one example:

                Table 4.2.5: 3-bit Complex Configuration Example #1
                BasebandSubbandPol'n productsSpectral channelsnBlBPQuadrant(s): Column(s)
                A1/C1 sb0-8 RR, LL, RL, LR 9 x 64 9 x 1 Q1: 0–8
                A1/C1 sb9 RR, LL 1 x 1152 1 x 9 Q1 & Q3: 9–11, 14 / Q4: 9
                A1/C1 sb10 RR 1 x 1792 1 x 7 Q1 & Q3 & Q4 : 12,13 / Q2: 13
                A1/C1 sb11 RR, LL 1 x 384 1 x 3 Q1 & Q2 & Q3: 15
                A2/C2 sb0-11 RR, LL, RL, LR 12 x 64 12 x 1 Q2: 0–11
                A2/C2 sb12 LL 1 x 768 1 x 3 Q2: 12, 14 / Q4: 14
                B1/D1 sb0-3 RR, LL, RL, LR 4 x 64 4 x 1 Q3: 0–3
                B1/D1 sb4 RR, LL, RL, LR 1 x 320 1 x 5 Q3: 4–8
                B2/D2 sb0-6 RR, LL, RL, LR 7 x 64 7 x 1 Q4: 0–6
                B2/D2 sb7 RR, LL 1 x 640 1 x 5 Q4: 7, 8, 10, 11, 15
                RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

                 

                Once again, the GOST implements all of these constraints, and is generally smart enough to figure out the routing scheme that works best for your particular request.

                Documentation

                Documentation

                Documentation for VLA data reduction, image making, observing preparation, etc., can be found in various manuals. Current manuals are available on-line. Those manuals marked by an asterisk (*) can be mailed out upon request, or are available for downloading from the NRAO website. Direct your requests for mailed hardcopy to Lori Appel. Many other documents of interest to the VLA user, not listed here, are available from our website.

                1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: "Synthesis Imaging in Radio Astronomy." This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series. The lectures of the 2014 workshop are available at the 14th Synthesis Imaging Workshop web site.
                2. PROCEEDINGS FROM THE 1998 SYNTHESIS IMAGING WORKSHOP: This is an updated and expanded version of Reference 1, taken from the 1998 Synthesis Imaging Summer School, held in Socorro in June, 1998. These proceedings are published as Volume 180 of the ASP Conference Series.
                3. GUIDE TO OBSERVING WITH THE VLA: Describes details of how to observe with the VLA once you have been allocated time on the VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide). Including special observing modes such as:
                  1. CALIBRATION (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/calibration)
                  2. OBSERVING WITH THE 8-BIT (up to 2 GHz bandwidth) & 3-BIT (up to 8 GHz bandwidth) SAMPLER SYSTEMS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/set-up);
                  3. SPECTRAL LINE OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/line); 
                  4. HIGH FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/hifreq);
                  5. LOW FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/lofreq);
                  6. VERY LOW FREQUENCY OBSERVING (< 500 MHz) (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/topical-guides/vlofreq);
                  7. POLARIMETRY (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/pol);
                  8. MOSAICKING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/mosaicking);
                  9. RADIO FREQUENCY INTERFERENCE (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/rfi);
                  10. MOVING OBJECTS (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/moving);
                  11. VLBI AT THE VLA (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/vlbi).
                4. *CASA COOKBOOK: The CASA Cookbook for use of the package for data reduction of VLA (& ALMA) data is available, along with other documentation, from the CASA home page (http://casa.nrao.edu). See (http://casa.nrao.edu/docs/cookbook/)
                5. VLA CASA Guides: Tutorials and data reduction examples of VLA data in CASA (https://casaguides.nrao.edu/index.php/Karl_G._Jansky_VLA_Tutorials)
                6. *AIPS COOKBOOK: The Astronomical Image Processing System (AIPS) software is able to fully calibrate VLA data and do most imaging operations. The exception is the wide-band (bandwidth synthesis) deconvolution which is being developed in CASA only. ALMA data may also be reduced in AIPS although the package is not fully qualified to calibrate data from the ALMA linearly-polarized feeds. The Cookbook description for calibration and imaging under the AIPS system can be found near all public workstations in the SOC. The latest version has expanded descriptions of data calibration imaging, cleaning, self-calibration, spectral line reduction, and VLBI reductions. See (http://www.aips.nrao.edu/cook.html)
                7. *GOING AIPS: This is a two-volume programmers manual for those wishing to write programs under AIPS. It is now somewhat out of date. See (http://www.aips.nrao.edu/goaips.html)
                8. *VLA CALIBRATOR LIST: This page contains the list of VLA Calibrators in both 1950 and J2000 epoch. See (https://science.nrao.edu/facilities/vla/observing/callist)
                9. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983.
                10. *HISTORICAL VLA MEMO SERIES: archive memo series from the early days of the VLA. See (http://library.nrao.edu/vlam.shtml)
                11. *RECENT VLA MEMO SERIES: the memo series relating to the expanded capabilities of the VLA. See (http://library.nrao.edu/evla.shtml)
                12. *The VLA Expansion Project: Construction Project Book. The Expanded VLA Project Books contains the technical details of the VLA Expansion construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml.
                13. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains a useful source of information on much of the VLA. There are a few hard copies at the VLA and in the DSOC. Much of this document is now available for download (https://science.nrao.edu/facilities/vla/obsolete/green-book). Note: it does not include any information about the hardware and software specific to the expanded Karl G. Jansky VLA.

                 

                Online Tools & Important Links

                The NRAO User Portal. (https://my.nrao.edu) This is a gateway to the NRAO interactive services that include the Proposal Submission Tool (PST).

                The NRAO Proposal Submission Tool (PST) online manual. (https://science.nrao.edu/facilities/vla/docs/manuals/proposal-guide/pst)

                The VLA Exposure Calculator Tool (ECT) online manual. (go.nrao.edu/ect)

                The VLA Exposure Calculator Tool (ECT). (https://obs.vla.nrao.edu/ect)

                The VLA GO Setup Tool (GOST). (https://science.nrao.edu/facilities/vla/docs/manuals/propvla/gost)

                 

                Acknowledgements

                Many thanks to all the VLA staff and our RSRO participants who have worked long and hard to commission these capabilities and who have helped to create this extensively updated set of documentation.

                NRAO is grateful to Professor Rob Ivison for supporting the upgrade of some of the 3-bit samplers on the VLA via a grant from the European Research Council. For observations using the 3-bit samplers between May 2015 and March 2018 we encourage users to include the following text in the Acknowledgments section of their publications:

                "We acknowledge funding towards the 3-bit samplers used in this work from ERC Advanced Grant 321302, COSMICISM."

                Contact Information

                Please go to the People page for information on key personnel at NRAO Socorro.

                Please direct queries to the NRAO Helpdesk; you can expect a response within one business day..  


                 

                Editor's Notes

                This Observational Status Summary for the Karl G. Jansky (expanded) VLA is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. For questions on the content, or suggestions that would enhance the clarity of this guide, we recommend contacting the NRAO Helpdesk.

                WIDAR Correlator: Spectral Line Observing (COPY)

                Basebands and Subbands

                Observers have access to very flexible correlator configurations using up to 64 subbands in up to 4 basebands sampled with the 8-bit and/or the 3-bit samplers. These capabilities may be summarized as follows:

                • Two 1 GHz baseband pairs using the 8-bit samplers, or four 2 GHz baseband pairs using the 3-bit samplers, independently tunable within the limits outlined in the section on VLA Frequency Bands and Tunability. The 8-bit baseband pairs are referred to as A0/C0 and B0/D0, while the 3-bit samplers are A1/C1, A2/C2, B1/D1, and B2/D2. The AC/BD nomenclature corresponds to that of the IF pairs in the pre-expansion VLA.
                • Up to 16 subband pairs (spectral windows) in each 3-bit baseband pair, and up to 32 subbands in each 8-bit baseband pair, for a total of up to 64 subbands in any correlator configuration:
                  • Tuning, bandwidth, number of polarization products, and number of channels can be selected independently for each subband;
                  • All subbands must share the same integration time;
                  • No part of a subband can cross a 128 MHz boundary;
                  • Subband bandwidths can be 128, 64, 32, …, 0.03125 MHz (128 / 2n, n=0, 1, …, 12).
                • The sum over subbands of channels times polarization products is limited to 16384 (without recirculation):
                  • These may be spread flexibly over subbands and polarization products, in multiples of 64: 64, 128, 192, 256, 384, …, 16384 cross-correlation products;
                  • Recirculation may be used to increase the number of channels per subband for subbands narrower than 128 MHz. Baseline Board stacking may be used to increase the number of channels per subband for setups requiring less than 64 subbands;
                  • Assigning many channels to a given subband may reduce the total bandwidth and/or the total number of subbands available.

                The remainder of this section discusses the various limitations in more detail, including some examples to show how they come up in practice.

                 

                Subband Tuning Restrictions

                Each subband may be placed anywhere within a baseband, with the caveat that no subband may cross a 128 MHz boundary. Mathematically:

                νBB0 + n×128 MHz <= νsbLow <= νsbHigh <= νBB0 + (n+1)×128 MHz

                where:

                νBB0 the lower frequency edge of the baseband;
                n= 0, 1, …, 7 (, …, 15) (any integer between 0 and 7 for 8-bit, between 0 and 15 for 3-bit);
                νsbLow                                   
                the lower edge of the subband
                (the subband center frequency minus half the subband bandwidth);
                νsbHigh                                 
                the upper edge of the subband
                (the subband center frequency plus half the subband bandwidth).

                For example, if the baseband were tuned to cover 10000–11024 MHz, one could place a 64 MHz subband to cover 10570–10634 MHz, but not to cover 10600–10664 MHz because that would cross the 128 MHz boundary at 10640 MHz. Note in particular that the center of a baseband is a boundary and no line should be observed at the baseband center.

                The figure below illustrates these restrictions:

                Correlator configuration figure: bandpass8jul12.png

                The black curve shows the analog filter response for an 8-bit baseband covering 1024 MHz, starting at νBB0. The dashed blue vertical lines show the 128 MHz boundaries; no subband can cross those boundaries and 128 MHz subbands are thus constrained to cover a region between two of those boundaries, with no finer tuning being possible. Narrower subbands, like the 64 MHz subband shown here in red, can be shifted around arbitrarily within one of the 128 MHz slots, but cannot cross any of these boundaries. (The dotted vertical red lines show the boundaries of the 64 MHz subband, while the solid curve shows an illustrative line within the subband.)

                The analog filter shape defining the baseband rolls off severely at one edge of the baseband, so the 128 MHz slot at that edge has reduced sensitivity. The baseband edge is at the lowest sky frequency in the baseband when using the upper sideband, and at the highest sky frequency in the baseband when using the lower sideband.

                 

                Subband Bandwidths and the Digital Filter Response

                The bandwidth for each subband may be selected independently, and can be any of 128/2n MHz, for n= 0, 1, …, 12: 128, 64, 32, 16, 8, 4, 2, or 1 MHz, or 500, 250, 125, 62.5, or 31.25 kHz.

                The usable portion of the subband is set by three effects. First, as discussed above, the analog filters which define the baseband are not perfect, leading to lower sensitivity in the 128 MHz near the baseband edge for the 8-bit samplers.

                Second, because the digital filters are not infinitely sharp, the rejected sideband leaks in at both edges of the subband. This leads to additional (aliased) noise, with a factor ~2 increase in the noise at the subband edges, dropping to a few percent within a few percent of the subband edge. The precise filter shape and noise increase is a complex but predictable function of the subband bandwidth (sbBW) and the subband tuning.

                The third effect stems from the offset frequencies used for sideband rejection in the WIDAR correlator. The local oscillators at the individual antennas are tuned to slightly different frequencies, with those offsets taken out in the correlator. This means that each antenna observes a slightly different sky frequency, and thus some baselines will not give an interesting correlation near one edge of the subband. The maximum frequency shift is currently set to 32×f0, with the fundamental f0 being set to f0 = max(25.6 kHz×sbBW/128 MHz, 100 Hz). Here sbBW is the smallest subband bandwidth within the baseband. For the wider subband bandwidths the maximum frequency shift corresponds to <1% of that bandwidth, but for narrower subbands the effect can be severe. For instance, a 31.25 kHz subband has f0 = 100 Hz, and a maximum frequency shift of 3.2 kHz—10% of the subband may be lost on some baselines.

                 

                Spectral Channels and Polarization Products

                Each subband, without recirculation enabled, can have a different number of channels and polarization products, subject to two limitations:

                1. For the ithsubband, the number of spectral channels can be:
                  • 64 nBlBP,i with full polarization products (RR,RL,LR,LL)
                  • 128 nBlBP,i with dual polarization products (RR and LL)
                  • 256 nBlBP,i with a single polarization product (RR or LL)
                  Here nBlBP,i= 1, 2, 3, 4, 5, …, 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.
                2. The sum over all subbands of nBlBP,i must be less than or equal to 64, the number of Baseline Board pairs in the correlator. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64 × 256 = 16,384 (without recirculation).

                Baseline Boards are the boards in the WIDAR correlator where the actual cross-multiplications are done. There are 128 Baseline Boards arranged as 64 Baseline Board pairs (BlBPs). The limitations given here correspond to the capabilities of the individual boards and the finite number of boards the correlator has.

                Limitation #1 corresponds to table 5.1.1 of the options for subband bandwidth and spectral resolution when using nBlBP Baseline Board pairs for a subband:

                Table 5.1.1: Subband Bandwidth and Spectral Resolution Options (without recirculation)
                Subband bandwidth &
                total velocity coverage
                Full polarization products
                (RR, RL, LR, LL)
                64nBlBP spectral channels

                Channel spacing:
                Dual polarization products
                (RR and LL)
                128nBlBP spectral channels
                Channel spacing:
                Single polarization product
                (RR or LL)
                256nBlBP spectral channels

                Channel spacing:
                128 MHz 38400/νGHz km/s 2000/nBlBP kHz 600/nBlBPGHz km/s 1000/nBlBP kHz 300/nBlBPGHz km/s 500/nBlBP kHz 150/nBlBPGHz km/s
                64 19200 1000 / nBlBP 300 / nBlBP 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP
                32 9600 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP
                16 4800 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP
                8 2400 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP
                4 1200 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 /n BlBP
                2 600 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP
                1 300 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP
                0.5 150 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP
                0.25 75 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP
                0.125 37.5 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP
                0.0625 18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
                0.0325 9.375 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP 0.122 / nBlBP 0.037 / nBlBP

                Subband bandwidth and spectral resolution options. Note that the table entries refer to the spacing between spectral channels—that spacing is before any frequency smoothing, so these channels are not independent.

                • nBlBP is the number of Baseline Board Pairs assigned to the subband.
                • Each subband may have a different number of spectral channels and polarization products, and each may be tuned independently.
                • There can be at most 16 subbands per baseband, and nBlBP must be an integer: 1, 2, 3, 4, 5, …, 64.
                • The sum of nBlBP over all subbands must be less than or equal to 64.
                • Use of more than one BlBP for a subband may further restrict the number of subbands available in one or more of the basebands; see text for details.

                Here are four examples of allowed general observing setups which use all 64 BlBPs to produce the maximum number of channels times polarization products:

                Table 5.1.2: Example BlBP Setups
                Baseband Subband

                Pol'n
                Products

                Spectral
                channels

                nBlBP
                Example 1 A0/C0 sb0 RR 16384 64
                Example 2 A0/C0 sb0 RR 8192 32
                A0/C0 sb1 RR, LL 1024 8
                A0/C0 sb2 RR, LL 512 4
                B0/D0 sb0 RR, LL 2048 16
                B0/D0 sb1 RR,RL,LR,LL 256 4
                Example 3 A0/C0 sb0 RR 8192 32
                A0/C0 sb1 LL 1024 4
                A0/C0 sb2 RR, LL 1024 8
                A0/C0 sb3 RR,RL,LR,LL 1024 16
                A0/C0 sb4 RR,RL,LR,LL 256 4
                Example 4 A0/C0 sb0-5 RR,RL,LR,LL 64 6 x 1
                A0/C0 sb6 RR, LL 3840 1 x 30
                A0/C0 sb7 RR 768 1 x 3
                A0/C0 sb8 RR,RL,LR,LL 192 1 x 3
                B0/D0 sb0-2 RR,RL,LR,LL 64 3 x 1
                B0/D0 sb3 LL 768 1 x 3
                B0/D0 sb4 RR, LL 2048 1 x 16

                 

                Recirculation

                Recirculation is a term to describe the method of increasing the number of spectral channels in a subband using correlator software, as opposed to Baseline Board stacking, which uses correlator hardware (see below). Recirculation is presently achieved by limiting the subband bandwidth and is available only for subbands less than 128 MHz wide. When limiting the bandwidth in a subband, the correlator software can be directed to use the remaining CPU cycles on a Baseline Board pair to obtain more lags in factors of two, running the data through the board for a second, third, etc., time; hence recirculation.

                At some time in the future an alternative method of recirculation, using the extra CPU cycles freed up by increasing the integration time, will be made available. Recirculation, by limiting of the subband bandwidth to increase the number of channels in factors of two, was used in the pre-expansion VLA correlator.

                 

                Recirculation vs. Baseline Board Stacking

                When faced with the choice between recirculation and Baseline Board stacking to increase the number of channels in a subband, we recommend recirculation for subbands narrower than 128 MHz and is supported in observatory software (GOST, OPT). For subbands of 128 MHz, Baseline Board stacking should be utilized to increase the number of channels.

                The present implementation of recirculation is that, for each halving of the subband bandwidth, the number of channels in the subband may be doubled without having to trade off the use of other subbands. Because recirculation is achieved by limiting the subband bandwidth, it is not supported for 128 MHz subbands; therefore for 64 MHz subbands a factor 2 recirculation is supported, etc. The maximum recirculation factor for a subband is 128/(subband bandwidth in MHz) and, of course, subject to other configuration restrictions such as data rate.

                The juggling between the requested number of channels, subband bandwidth, and the available number of Baseline Board pairs is dependent on the science goals and not easily formulated in a standard answer. However, if subbands of less than 128 MHz are used, recirculation becomes an option for setups that can also be achieved with Baseline Board stacking. In such cases, we suggest to use recirculation where possible and within the General Observing or Shared Risk requirements. This frees up unused Baseline Board pairs for other use; alternatively, one becomes less dependent on all Baseline Board pairs being in working order.

                Recirculation with factors 8 to 64 is designated Shared Risk, and recirculation with factors over 64 as Resident Shared Risk. The latter choice may have severe implications for the sensitivity as visibility integration time is used as trade-off. Ask the NRAO Helpdesk for more details.

                 

                Baseline Board Stacking

                As opposed to recirculation, which increases the number of channels in a subband by exploiting otherwise unused CPU resources, Baseline Board stacking adds more channels to a subband by adding correlator hardware resources, i.e., using up more Baseline Board pairs. Using Baseline Board stacking may therefore limit the number of subbands available in one or more of the basebands. Understanding how this works requires understanding some of the details of the correlator hardware. That understanding is built into the GOST, and observers may simply use that tool to find out whether their particular setup will, in fact, work. But the results can be confusing without some understanding of the hardware constraints from which they arise. These hardware constraints are complex, and most observers will not need to understand these details. The following section is for those who are attempting complex line experiments and who find the GOST or the RCT restricting the number of subbands and/or channels they can use in unexpected ways. Most observers can skip the following section.

                Baseline Board Stacking and Correlator Use

                First let us consider how the correlator hardware is organized. The cross-multiplications in the WIDAR correlator are spread across 64 Baseline Board pairs (BlBP), arranged into 4 quadrants of 16 BlBP each. Each baseband is connected directly to one of those quadrants. In the simplest mode, each of the 16 BlBP of a quadrant handles the correlations for one of the 16 subbands of the corresponding baseband. Four basebands and four quadrants are required to handle the full 8 GHz of bandwidth per polarization provided by the 3-bit (wideband) samplers: 8 GHz is split into four basebands of 2 GHz each, with each baseband fed into a different BlBP quadrant. Each BlBP in that quadrant handles a subband of maximum bandwidth 128 MHz, so 16 BlBP handles 16 subbands for a total of 16×128 MHz = 2048 MHz.

                A single BlBP produces 256 cross-correlations per baseline for a single subband, which can be used for a single polarization product (e.g., RR or LL with 256 spectral channels), or two (RR and LL with 128 spectral channels each), or four (RR, RL, LR, and LL with 64 spectral channels each).

                 

                When using the 8-bit samplers, the total bandwidth is only 2 GHz per polarization, split into two basebands of 1 GHz each. The simplest continuum setup uses only two quadrants, since there are only two basebands; and only 8 subbands are required to span the 8×128 MHz = 1024 MHz of each baseband. Three-quarters of the correlator BlBP hardware remain unused.

                 

                The spectral line mode allows access to these extra correlator resources through Baseline Board stacking: using multiple BlBPs to process the same subband and produce more cross-correlations for that subband. This is done using crossbar switches which make the data for a single subband available to several BlBPs. Those BlBPs can then be used to produce more spectral channels for that subband, with n BlBPs producing 256×n cross-correlations per baseline. The limit on the total number of cross-correlations (16384) stems from the total number of BlBPs (64): 64×256 = 16384.

                Unfortunately, completely flexible crossbar switches are expensive and could not be implemented in the VLA's correlator. This means that one cannot route a given subband to a randomly-chosen BlBP. The routings which are possible, are as follows:

                1. A subband in a baseband can be routed to any BlBP within the corresponding quadrant.
                2. Data coming into a given BlBP in one quadrant, can be routed to the corresponding BlBP in any other quadrant.

                Routing option #1 means that one could use all the BlBPs within a quadrant to correlate a single subband, yielding 16×256 = 4096 cross-correlations for that subband:

                 

                Routing option #2 means that one could use the BlBPs in all 4 quadrants to correlate a single subband. One simple case would use 4 BlBPs to correlate each of the 16 subbands in a single baseband, yielding 4×256 = 1024 cross-correlations for each of those subbands. Note that in this case, no BlBPs are left to correlate any data from the second baseband.

                Using routing option #2 does come with a subtle cost: assigning a BlBP in quadrant X to correlate a subband corresponding to quadrant Y removes that BlBP from use in the baseband corresponding to quadrant X…and therefore also removes the corresponding subband in that baseband. So, getting more channels for a subband in one baseband may prevent the use of a subband in a different baseband. To take a simple example, consider an experiment where one wishes to observe a single line in dual polarization with 512 channels (requiring 4 BlBPs), plus as much continuum bandwidth as possible. Naively, one would say there are 16 subbands in each baseband; one is used for the spectral line, so that leaves 16+15 = 31 subbands, and with the widest subband bandwidth (128 MHz) the total available continuum should be 31×128 MHz = 3968 MHz per polarization. Actually, however, there are only 15+15 subbands available, or 30×128 MHz = 3840 MHz per polarization, because the spectral line subband has eaten one BlBP corresponding to the other baseband:

                 If the same spectral line required twice as many channels, this will result in the loss of two subbands in both of the basebands:

                In some cases one may want to use a different routing to use up subbands in one baseband in preference to another. For instance, the same spectral line setup (2048 cross-correlations for a single spectral line subband, plus as much continuum as possible) could be set up to allow 13 continuum subbands in the A0/C0 baseband, and the full 16 continuum subbands in B0/D0:

                Understanding these confusing constraints can help observers set up the VLA more effectively to achieve their scientific goals. For instance, in a mixed line+continuum experiment, it works best to use the resource tools to set up the baseband tunings and subband channelization for the most important lines first, then add the desired continuum, and then see what correlator resources remain for any lines of secondary interest.

                The above examples all use BlBP pair stacking in powers of 2, but this is not required. To give some idea of more complex possibilities, the following tables (5.1.3 and 5.1.4) give two examples of other possible configurations. The RCT display shows how the Baseline Boards are used to process the individual subbands. The cyan boxes (shaded when printed out in black and white) show the Baseline Boards used to process data from baseband A0/C0, while the yellow boxes show Baseline Boards used to process data from baseband B0/D0.

                Table 5.1.3: Complex Configuration Example #1
                BasebandSubbandPol'n productsSpectral channelsnBlBP
                A0/C0 sb0 RR 10240 40
                A0/C0 sb1 LL 768 3
                A0/C0 sb2 RR,LL 2176 17
                B0/D0 sb0 RR 256 1
                B0/D0 sb1 RR,LL 384 3
                RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

                 

                Table 5.1.4: Complex Configuration Example #2
                BasebandSubbandPol'n productsSpectral channelsnBlBP
                A0/C0 sb0 RR 4352 17
                A0/C0 sb1 RR, LL 1152 9
                B0/D0 sb0 RR,RL,LR,LL 192 3
                B0/D0 sb1 RR, LL 4480 35
                RCT display: corr-cfg-fig:sro2_8bit_ac17+9_bd3+35

                 

                The individual subbands can have different bandwidths, and those bandwidths may be chosen completely independently of the number of spectral channels in each subband. So, for instance, a subband with a bandwidth of 2 MHz and 1152 spectral channels would have a channel separation of 2 MHz/1152 = 1.736 kHz; but the observer could equally well choose a bandwidth of 64 MHz for that subband, leading to a channel separation of 64 MHz/1152 = 55.56 kHz.

                Use of the 3-bit samplers further extends the possibilities.  Here is one example:

                Table 5.1.5: 3-bit Complex Configuration Example #1
                BasebandSubbandPol'n productsSpectral channelsnBlBPQuadrant(s): Column(s)
                A1/C1 sb0-8 RR, LL, RL, LR 9 x 64 9 x 1 Q1: 0–8
                A1/C1 sb9 RR, LL 1 x 1152 1 x 9 Q1 & Q3: 9–11, 14 / Q4: 9
                A1/C1 sb10 RR 1 x 1792 1 x 7 Q1 & Q3 & Q4 : 12,13 / Q2: 13
                A1/C1 sb11 RR, LL 1 x 384 1 x 3 Q1 & Q2 & Q3: 15
                A2/C2 sb0-11 RR, LL, RL, LR 12 x 64 12 x 1 Q2: 0–11
                A2/C2 sb12 LL 1 x 768 1 x 3 Q2: 12, 14 / Q4: 14
                B1/D1 sb0-3 RR, LL, RL, LR 4 x 64 4 x 1 Q3: 0–3
                B1/D1 sb4 RR, LL, RL, LR 1 x 320 1 x 5 Q3: 4–8
                B2/D2 sb0-6 RR, LL, RL, LR 7 x 64 7 x 1 Q4: 0–6
                B2/D2 sb7 RR, LL 1 x 640 1 x 5 Q4: 7, 8, 10, 11, 15
                RCT display: corr-cfg-fig:sro1_8bit_ac40+3+17_bd1+3

                 

                Once again, the GOST implements all of these constraints, and is generally smart enough to figure out the routing scheme that works best for your particular request.

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                VLA OSS correlator configuration figure: 3bit_4x16x128MHz.png

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                VLA OSS correlator configuration figure: 3bit_4x16x128MHz.png

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