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VLA Observational Status Summary 2013B

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VLA capabilities September 2013 - January 2014

1. Introduction

1.1. Purpose of Document

The Observational Status Summary (OSS) summarizes the instrumental status of the Karl G. Jansky Very Large Array (VLA) for the CnB, B & BnA configurations starting August 23, 2013, and should be used when preparing for proposing for the February 1, 2013 deadline.  For capabilities prior to that date, please consult the previous version of the OSS.

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 consult the VLA's staff members, listed in Key Personnel, 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 Help for Visitors to the VLA and SOC for details. The VLA is an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it.

1.2. The Expanded Very Large Array (EVLA) project

The Expanded VLA (EVLA) project is a program to modernize the electronics of the Very Large Array (VLA) that was 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 on the web, at http://www.aoc.nrao.edu/evla/.   The EVLA project is 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 $94 million in Year 2006 dollars, including $59 million in new NSF funding, $16 million in redistributed effort from the NRAO Operations budget, $17 million for the correlator from Canada, and $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 that can process 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 will be an improvement over 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 who could be in 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.  The full 8 GHz bandwidth is now being offered to the general community at the start of full operations, semester 2013A.

1.3. VLA Science Opportunities

1.3.1. General

The 2013B call for proposals (http://science.nrao.edu/enews/6.1) details the general capabilities being offered for the Karl G. Jansky Very Large Array (VLA) in the 2013B semester (August 23, 2013 to January 6, 2014) with flexible tuning of sub-band spectral line windows using the 8-bit samplers (enabling up to 2 GHz total bandwidth), use of the 3-bit samplers at higher frequencies in a mode that is suitable for wide-band continuum and extragalactic lines and line searches, use of up to 3 independent sub-arrays and a phased array capability for VLBI.

In addition to the 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" and "Resident Shared Risk Observing" (RSRO) programs.     Details of what is being offered for the shared risk programs and what is required to write a proposal for them are available on the web at: https://science.nrao.edu/facilities/vla/proposing/sro2013b.

1.3.2. P-band

P-band (230-470 MHz) is available for shared-risk proposals in 2013b.

The new EVLA Lowband receiver system consists of two bands, P: (230-470 MHz) and 4: (54-86 MHz).  Lowband is in the final stages of construction and commissioning has begun. The 4-band part of the system is not ready for open observing yet but P-band is ready for simple, shared-risk continuum projects.  We expect at least 20 antennas to be working at P-band during 2013b. The Lowband system is mounted near the prime focus and, unlike the rest of the VLA bands, uses linear polarization. The wide-band system is significantly more sensitive than the old VLA P-band, due to a lower system temperature and the use of the WIDAR correlator.  We are not yet ready to provide polarization calibration but the system works adequately for simple continuum observations of unpolarized sources. For now the default path for calibration uses AIPS; however, for most projects, the imaging is best done in CASA using its wideband clean algorithm with nterms=2. Calibrated uv data can be ported from AIPS to CASA.

2. 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 relevant postscript file.

The VLA completes one cycle through all four configurations in an approximately 16 month period.    The VLA configuration schedule for 2013/2014 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/facilities/evla/proposing/configpropdeadlines.shtml or recent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines.  Refer to "VLA Proposal Preparation and Submission"  for information on how to submit an observing proposal.

Table 3: VLA Configuration Schedule for 2013/2014

25 Jan - 22 Apr, 2013 26 Apr - 13 May, 201324 May - 19 Aug, 2013 23 Aug - 9 Sep, 201320 Sep - 16 Dec, 2013 20 Dec - 6 Jan, 201417 Jan - 28 Apr, 2014
Configuration: D

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, the VLA is now being outfitted with even lower frequency receivers, P-band (230-470 MHz) and 4-band (54-86 MHz).  Up to 20 P-band receivers are expected to be fully operational in mid-2013 while the 4-band systems are still undergoing development at this time.

The VLA correlator is both powerful and flexible.  Details of the correlator configurations being offered for VLA science in 2013  are described in the Observational Status Summary - Correlator Configurations.   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.

3. Performance of the VLA

3.1. Introduction

This section contains details of the EVLA's resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Documentation].

3.2. 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/D) 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 fully recover this missing information, which can only be obtained by observing in a smaller array configuration, by using the mosaicing method, or by utilizing data from an instrument (such as a large single antenna or an array comprising smaller antennas) which provides this information.

Table 5 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.
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 will 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.

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; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed.

3.3. Sensitivity

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




- 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 (at least 0.92 for the VLA).
- 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 2 shows the SEFD as a function of frequency for those bands currently installed on VLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial VLA frequencies.


Note: SEFDs at Ku, K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions.

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 6 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 6; (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 = 0.8 to 0.85.
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 be 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.

    The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity, generated from ground-based weather, in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves.

  • 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 exposure tool calculator on-line at http://science.nrao.edu/facilities/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations.

Special caveats apply for P-band (235 -- 490 MHz) observing.  The listed SEFD in table 6 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:

[display]T_{\rm b} = \frac{S \lambda^2}{2k\Omega} = F \cdot S[/display]

...equation (2)

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 constant F depends only upon 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):

[display]M_{\rm HI} = 2.36 \times 10^5 D^2 \sum S \Delta V ~M_\odot[/display]

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


Also, look at Figure 1.

3.4. VLA Frequency Bands and Tunability


For observations taken with the 8-bit samplers under the general observing program,  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.  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, IFs C and D provide LCP.  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.   Currently, a maximum of 1.024 GHz can be correlated for each IF pair (see Correlator Configurations), for a total maximum bandwidth of  2.048 GHz.  To distinguish this 8-bit system from the 3-bit system (described below), these IF pairs are denoted A0/C0 and B0/D0.

The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below.

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 default setup for P-band will provide two subbands from the A0D0 IF pair, each 128 MHz wide, spanning 236 -- 364 MHz, and 364 -- 492 MHz.  The channel resolution is 125 kHz.   The B0D0 pair will be used for the 4-band tuning. 
3. 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.

In general, for all frequency bands except Ka, if the total span of the two independent IFs (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 AC IF must be greater than 32.0 GHz. There is no restriction on the BD frequency, unless the BD band overlaps the AC 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 BD frequency must be lower than the AC frequency.

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 A1C1 and A2C2, and the B/D IF pair provides two sampled pairs, labelled B1D1 and B2D2.  The maximum frequency span permitted for the A1C1 and A2C2 pairs is about 5000 MHz.  The same restriction applies to B1D1 and B2D2.   The tuning restrictions given above for the separation and location of the 8-bit pairs A0C0 and B0D0 also apply to the 3-bit pairs.

3.5. VLA Samplers

The VLA is now equipped with two different types of samplers.  Which set you should use depends primarily on your science goals.

A)  8-bit Set

This set consists of four 8-bit samplers running at 2048 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, but some restrictions apply, depending on band.  Tuning restrictions for these two pairs are described in the Frequency Bands and Tunability section of this document.

B)  3-bit Set

This set consists of eight 3-bit samplers running at 4096 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.  A more complete description of the tuning restrictions is given in the VLA Frequency Bands and Tunability section of this document.

C)  Major Characteristics of each Set

The 8-bit samplers have been in use for many years, and their characteristics are well understood.  Use of this sampler set is warranted for low-band, L-band, and S-band observations, where the full analog bandwidth provided by the receivers is less than or equal to the 2048 MHz span covered by the samplers.

The 3-bit samplers, at the time of writing, are still being implemented on the array.  Their characteristics are currently being explored, so users should refer to other sections of this document for updated and current instructions for optimum use.   Major issues users need to be aware of include:

  • The sensitivity of the array is reduced by 10 to 15% with the 3-bit system, compared to the ideal analog system.  For comparison, note that the loss is less than 3% for the 8-bit system.
  • Each of the eight 3-bit samplers has a resonance of approximate width 3 MHz.  Each sampler's resonance is independent of, and different from, all others, so there is no correlated signal between antennas.  The resonance, in effect, greatly reduces the SNR within the narrow frequency span of that resonance.  This will degrade bandpass solutions within the narrow resonance window, and images made at frequencies within any one of the resonances will show significant loss in sensitivity.  The resonance effect is easily seen in autocorrelation spectra, and it is recommended that users utilize these to inform themselves of the compromised frequencies.  As not all samplers are (as of this writing) outfitted on the array, we can not yet provide a table of the compromised frequencies.
  • The SY (switched power) values utilized to correct for system gain variations are sensitive to the total power from each antenna.   Application of the switched power values will bias the resulting visibilities by a value of 5 to 10%.  This origin of this effect is well understood, but we have not yet determined how best to compensate for it.  Because of this, we do not recommend use of the SY table data for data taken with the 3-bit samplers.

D) Setting up the 3-bit or 8-bit Samplers

Setting up either set of samplers requires a short initial observation for each individual LO (frequency) tuning.  For the 8-bit system, a scan of 1 minute duration is sufficient for each tuning.  The pointing direction of the antennas is not critical.

For the 3-bit system, the requirements are more demanding.  The recommended duration is two minutes per frequency tuning.  In addition, we recommend that the antennas be pointed at an elevation near to 30 degrees.  A slew time of up to three minutes will likely be required to position the antennas at this elevation.

E)  Which set to use?

Observations at S, L, and low bands should use the 8-bit sampler set in all cases.  There is no benefit from use of the 3-bit samplers at these bands.

Continuum-application observations at Ku, K, Ka, and Q bands, where the best sensitivity is critical, should use the 3-bit sampler sets.  Similarly, wide-band spectral line searches requiring more than the 2 GHz span of the 8-bit set should use the 3-bit sampler sets.

Spectral-line application observations whose frequency span lies within the two 1024 MHz-wide spans offered by the 8-bit set should use the 8-bit set.

It is not known at this time if full-band continuum observations at C and X bands, where the maximum bandwidth is 4 GHz, should use the 8-bit or 3-bit sets.  In principle, the 3-bit system should offer ~40% better continuum sensitivity, but this will be eroded by:

  • the reduction of 10 to 15% noted above,
  • the increased setup time needed for 3-bit observations,
  • the extra RFI present at these bands which can be avoided by tuning the 8-bit system.  This is especially notable at C-band.

F) Other issues to be aware of.

The 8-bit system provides much better protection against strong RFI.  However, we do not yet know if the 3-bit system will be
significantly compromised by RFI at any band where its use is warranted.  This is most likely to be an issue at C-band, where
strong RFI from broadcast satellites and local microwave links are known to be present.  Our current understanding of the RFI spectrum at higher frequency bands suggests there will be no problems with the 3-bit system, except possibly when pointing near geostationary satellites.

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

3.6. Field of View

3.6.1. Introduction

At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly:

3.6.2. 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

3.6.3. 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 8 shows the decrease in peak response and the increase in apparent radial width as a function of this parameter. Table 8 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.

3.6.4. 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 9, 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.

3.6.5. 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 (http://www.aoc.nrao.edu/evla/geninfo/memoseries/evlamemo67.pdf) for more details.

3.7. Time Resolution and Data Rates

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

Configuration Observing
integration time
D, C
5 seconds
D, C
X Ku K Ka Q
3 seconds
B all 3 seconds
A all 1 second

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
Standard 1 sec
20 MB/s ( ~70 GB/hr)
Shared risk 50 msec 60 MB/s (~210 GB/hr)

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 × (Nant1)/(27×26) / (Δt/1 sec)
= 160 GB/hr × (Nchpol/16384) x Nant × (Nant1)/(27×26) / (Δt/1 sec)
= 3.7 TB/day × (Nchpol/16384) x Nant × (Nant1)/(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 ~3 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.

3.8. 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 10 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 online, at https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/rfi. 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 new 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 (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.  At C-band however, there are strong TV and microwave communications signals which may saturate the 3-bit samplers.  Until VLA staff can better measure this effect, we recommend use of the 8-bit samplers at C-band.

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.

3.9. Subarrays

The separation of the VLA into up to three subarrays will be supported for general observing in semester 2013B.  Each subarray is independent for the entire duration of an observation, and is set up using an independent scheduling block (SB).  All SBs must be the same length, and will be scheduled simultaneously.  The antennas comprising each subarray will remain in the same subarray for the duration of the observation.  There are some limitations on the number of antennas that can comprise each subarray, due to the architecture of the correlator.  The correlator operates on 8 groups of 4 antennas (up to 32 antennas total), and each group can only contain antennas from one subarray.  Thus, for three subarrays and 27 antennas, antennas may be distributed as 8, 9, 10 antennas, or 11, 11, 5 antennas, but not 9, 9, 9 antennas since each subarray of 9 antennas requires 3 groups, and the three subarrays together would require 3x3=9 groups, where only eight are available.  For two subarrays the grouping requirement gives no restrictions, and the antennas may be flexibly allocated between the two subarrays.

Subarrays will only be available for general observing with the 8-bit samplers and 2 GHz bandwidth.  Access to the 3-bit samplers and/or more flexible correlator set-ups will be offered under the shared risk observing programs.

3.10. Positional Accuracy

Summary: A target's position 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 of a target 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 phases as determined on the reference source. (Note that phase self-calibration imposes the assumed position of the model, i.e.,  makes  the position indeterminate.)

It may help to think of astrometry in 2 steps, narrow and wide-field. In the former, 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. Under more typical conditions, 10% of the beam is readily achieved.

Astrometric calibrators are marked '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.  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.

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.

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. These effects are handled somewhat differently in the various reduction packages. 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.

3.11. Limitations on Imaging Performance

3.11.1. 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 atmospheric phase stability, although pointing errors and changes in atmospheric opacity can also be a limiting factor.  If the target source contains compact structures of sufficient strength (depending on the band, bandwidth, atmospheric coherence time, and source complexity), self-calibration can be counted on to improve the images.  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.

3.11.2. 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 5 gives the largest scale visible to each configuration/band combination.

3.11.3. 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.

3.11.4. Sidelobes from Confusing 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.

3.11.5. 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.

3.11.6. Wide-band Imaging

The very wide bandpasses provided by the Jansky 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.

3.11.7. 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.

3.12. Calibration and Flux Density Scale

The VLA Calibrator List contains information on 1860 sources sufficiently unresolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/.

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, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well.  For S-band, use the L or C band models.

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 11 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 log2(f) + 0.0336 log3(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.

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 approximately every other year, 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.

Although the VLA's new electronics are very stable, gain variations in the system will occur, either due to changes in amplifier gains, or due to changes in internal attenuator settings made by the observing system to ensure that correct voltage levels are provided to the samplers.  These changes are monitored by an internal calibration signal, the results of which are stored in the switched power table (SY table).  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 SY table data.  For the 3-bit system there is an additional complication, as the values of the SY 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 SY table data will reduce bootstrapping accuracy to perhaps 5 to 10%.

3.13. Complex Gain Calibration

3.13.1. 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 impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below.

  • Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm.
  • 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 - once 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 newly-enhanced sensitivity of the VLA now guarantees, for full-band continuum observations, that every field will have enough background sources to enable phase self-calibration at L and S bands, and probably also at C-band.  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 (see http://www.vla.nrao.edu/astro/calib/manual/key.html for a description of calibrator codes).
  • 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 Figure 1 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


Figure 1 - 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.

3.13.2. 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. 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 EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Documentation.

Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells - most notably, thunderstorms).

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 can be used to estimate the required calibration cycle times when using fast switching phase calibration, and in the worst case, to indicate to the observer that high frequency observing may not be possible with current weather conditions.

Plots of current/historical data can be found at: https://webtest.aoc.nrao.edu/cgi-bin/thunter/api.cgi

Characteristic seasonal averages are represented below:

Click on the Month links above to see plots of phase and wind speed versus time.

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

3.14. 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.


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


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. In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5-5.0 GHz frequency range for C band, and outside 1.3-1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements.

3.15. Correlator Configurations

3.15.1. General Capabilities Introduction

The correlator configurations offered for general observing in the 2013B semester 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 proposals for the 2013B semester. We are in the process of updating the Resource Configuration Tool (RCT) to reflect the new possibilities.

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 setup provides the widest possible bandwidth for a given observing band, with channel spacing depending on the number of polarization products.

Wideband & Subarray Correlator Options
Polarization products Channel spacing
Full (RR, RL, LR, LL) 2   MHz
Dual (RR, LL) 1   MHz
Single (RR or LL) 0.5 MHz

For the low-frequency bands (L, S, C, X, and Ku bands) this uses the 8-bit samplers to provide a total of 2 GHz of bandwidth per polarization. The three highest-frequency bands (K, Ka, and Q bands) employ the new 3-bit (wideband) samplers, providing a total of 8 GHz of bandwidth; in this case all subbands are 128 MHz wide. 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. The (fairly mild) tuning restrictions are described in the section on VLA Frequency Bands and Tunability. Spectral Line Configurations

In semester 2013B observers have access to very flexible subbands within two 1 GHz baseband pairs (using the 8-bit samplers). These capabilities may be summarized as follows:

  • Two 1 GHz baseband pairs (using the 8-bit samplers), independently tunable within the limits outlined in the section on VLA Frequency Bands and Tunability. These baseband pairs are referred to as A0/C0 and B0/D0, and correspond to the IF pairs in the pre-expansion VLA.
  • Up to 16 subband pairs (spectral windows) in each baseband pair
    • 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.
    • These may be spread flexibly over subbands and polarization products, in powers of two: 64, 128, 256, 512, ..., 16384 cross-correlation products.
    • Assigning many channels to a given subband may reduce 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


νBB0 the lower frequency edge of the baseband;
n= 0, 1, ..., 7 (i.e., any integer between 0 and 7);
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).

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 can have a different number of channels and polarization products, subject to two limitations:

  1. For the ith subband, 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, 4, 8, ..., 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.  [Shared risk observations are allowed to set nBlBP,i to any integer between 1 and 64 -- see the section on Shared Risk Observing.]
  2. The sum over all subbands of nBlBP,i must be less than or equal to 64. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64x256= 16,384.

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
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
18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
0.03.25 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.
  • For general observations there can be at most 16 subbands per baseband, and nBlBP must be a power of two: 1, 2, 4, 8, 16, 32, or 64.
  • For shared risk observations nBlBP can be any integer between 1 and 64. Further, up to 32 subbands per 8-bit baseband are available.
  • In all cases 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





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 sb0-7 RR,RL,LR,LL 64 8 x 1
A0/C0 sb8 RR, LL 4096 1 x 32
B0/D0 sb0-7 RR,RL,LR,LL 64 8 x 1
B0/D0 sb8 RR, LL 2048 1 x 16

Number of subbands and Baseline Board stacking

The number of subbands available in a baseband depends on the correlator resources available.  This means that using Baseline Board stacking to increase the number of channels in one subband, may 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.  This 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.

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 Shared Risk Observing (SRO) programs provide even more options (see the ones for SRO and RSRO) and hence benefit even more from a clear understanding of the correlator hardware.

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. Continuum Subarrays

The continuum subarray option offers two 1 GHz baseband pairs (two 8-bit baseband pairs, each with 1 GHz in RCP and 1 GHz in LCP) 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 column Number
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

3.15.2. Shared Risk Observing Program

The Shared Risk Observing (SRO) program allows observers access to capabilities that can be set up via the Observation Preparation Tool and run through the dynamic scheduler (without intervention), but are not well tested. In semester 2013B this program provides a number of extended correlator capabilities:

  1. Data rates as high as 60 MB/s and integration times as short as 50msec.
  2. Use of the 3-bit samplers in any frequency band, with subbands which can be narrower than 128 MHz. All subbands within a baseband must be contiguous, and must have the same bandwidth.  Note that it makes little sense to use the 3-bit samplers for L and S bands, where the 8-bit system cover the entire available observing band with higher dynamic range and better sensitivity.
  3. Mixing 3- and 8-bit basebands (3-bit AC and 8-bit BD, or vice versa) in the same correlator configuration.
  4. More flexible Baseline Board stacking, with the ability to assign any number of Baseline Board pairs (BlBPs) to each subband, rather than only powers of two. These configurations are significantly more complex, and the observer is urged to read the section on spectral line capabilities (above) with care. Some examples of possible shared risk configurations are given below.
  5. Access to the 3-bit samplers for observations with (up to 3) subarrays.

Those proposing for Shared Risk Observing must explicitly define their correlator set-up(s) in the PST, instead of using the GOST. Those in doubt as to whether their proposed setups are Shared Risk or Resident Shared Risk are encouraged to get in touch with NRAO staff well before the proposal deadline.


Shared risk observations: spectral line configurations

Arbitrary Baseline Board stacking using the 8-bit samplers is allowed under the Shared Risk Observing (SRO) program for semester 2013B. The basic capabilities are given in the table above (Subband Bandwidth and Spectral Resolution Options), the difference being that the number of Baseline Board pairs assigned to a subband (nBlBP) can be any integer up to 64, rather than only a power of two. Similarly the interaction between BlBP stacking the number of independent subbands is as described above, but this interaction is more subtle and potentially confusing in the shared risk case.

To give some idea of the possibilities, the following tables give two examples of possible shared risk configurations. The SRCT 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.

Shared Risk Configuration Example #1
BasebandSubbandPol'n products Spectral 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
SRCT display: Correlator configuration figure: sro1_8bit_ac40+3+17_bd1+3.png




Shared Risk Configuration Example #2
BasebandSubbandPol'n products Spectral 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
SRCT display: Correlator configuration figure: sro2_8bit_ac17+9_bd3+35.png


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.

Again, can be difficult to distinguish between configurations which are available as Shared Risk Observing (SRO), and those which are only available through the Resident Shared Risk Observing (RSRO) program. If in doubt, please consult with NRAO staff.

3.15.3. Resident Shared Risk Observing Program

Even more sophisticated correlator configurations are available under the Resident Shared Risk Observing (RSRO) program. Examples include:

  • More extensive use of recirculation to provide many times more spectral channels, flexibly distributed among the subbands.
  • Access to more than 16 subbands per 8-bit baseband.
  • Data rates beyond 60 MB/s and integration times below 50 msec.
  • Pulsar phase bins and real-time de-dispersion.
  • Multiple phase/delay centers.
  • Access to wideband phased array output for high time resolution observations.
  • Real-time RFI excision and transient detection.
  • 7-bit correlations.
  • More subarrays, and more complex correlator configurations for each subarray.
  • Multiple integration times within a single subarray (e.g., slower integrations for spectral line than for continuum subbands).

Those interested in these capabilities should consult with NRAO staff, and be prepared to make a significant commitment to aid in commissioning the desired mode(s).

Correlator configurations and observing details for RSRO programs must be fully explained and justified in the Scientific and Technical Justification section of the proposal.

3.16. VLBI Observations

The VLA can now participate in VLBI observations with the VLBA.  Currently, this is only allowed in phased array mode (i.e. single dish is not allowed) with highly restricted WIDAR correlator configurations.   For more details see the VLBI at the VLA documentation.  In phased array mode the program TelCal derives the antenna-based delay and/or phase corrections needed for antenna phasing in real time.  This delay correction is then applied to the antenna signals in the correlator and the visibilities are then summed.  The resulting phased array sum is recorded in VDIF format on the Mark5C recorder 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.  The standard WIDAR data are also archived in the NRAO science data archive.

3.17. 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 in the D configuration, for which 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. Use of the AIPS program 'IMAGR' or CASA task 'clean' 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 "Co-Ordinated Radio 'N' Infrared Survey for High-mass star formation (CORNISH, http://www.ast.leeds.ac.uk/Cornish/public/index.php)" (B configuration), or the NRAO VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) (D configuration, all-sky) surveys.

3.18. 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 task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag modern VLA data based on shadowing as well, although it will only flag based on antennas in the dataset, and is ignorant of antennas in other sub-arrays.

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.

3.19. Combining Configurations and Mosaicing

Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. 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 EVLA are given in the VLA Observing Guide.

3.20. Pulsar Observing

While the WIDAR correlator will eventually allow flexible pulsar phase-binning and real-time correction for dispersion, those aspects of the correlator have not yet been commissioned.  Similarly the use of the phased array for high time-resolution observations is not yet possible for general use.

4. Using the VLA

4.1. Obtaining Observing Time on the VLA

Observing time on the VLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the VLA is based upon the submission of a VLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information.

It is also possible to obtain VLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information.

Students planning to use the VLA for their PhD dissertation must submit a "Plan of Dissertation Research" of no more than 1000 words with their first proposal. This plan can be referred to in later proposals. At a minimum the plan should contain a thesis time line and an estimate of the level of VLA resources needed. The plan provides some assurance against a dissertation being impaired by an adverse review of a proposal when the full scope of the thesis is not seen. The plan can be submitted via NRAO Interactive Services. Students are reminded to submit their plan comfortably in advance of the proposal deadline. New thesis plans must be in pdf format so science reviewers can easily access the plans. Students who have not yet graduated but have active plans on file should consider updating those plans to a pdf format if they are not already in that form.

Time on the VLA is scheduled on a semester basis, with each semester lasting six months. Proposal deadlines are 5pm (1700) Eastern Time on February 1 and August 1. If the deadline falls on a holiday or weekend, it is extended to the next working day. The February 1 proposal deadline nominally covers time to be scheduled during the following August through January, and the August 1 deadline is for time to be scheduled from February through July. At either deadline requests can be made for a VLA configuration linked to a future deadline.

Proposals are evaluated on the basis of scientific merit by eight Science Review Panels (SRPs). The proposals are also reviewed for technical feasibility by NRAO staff. The SRP's comments and linear-rank scores are strongly advisory to the Time Allocation Committee (TAC). The scores and comments from both the SRPs and TAC are passed on to the proposers soon after each meeting of the TAC (twice yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/observing/proposal-types/peta for a detailed description of the proposal evaluation and time allocation process.

Because of competition, even highly ranked proposals are not guaranteed to receive observing time. This is particularly true for proposals that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo Cluster. Daytime observing will also continue to be limited by VLA enhancement activities throughout 2013.

4.2. Director's Discretionary Time

The NRAO has established two categories of proposals for Director's Discretionary Time (DDT). DDT is limited to a maximum of about 5% of the total observing time on the VLA. All DDT proposals must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered.

1. DDT Target of Opportunity. Target of Opportunity (ToO) proposals are for unexpected or unpredicted phenomena such as extreme X-ray or radio flares. ToO proposals are evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. ToO proposals are evaluated on the basis of scientific merit by Observatory scientific staff. Those staff also asses the technical feasibility of the ToO proposals. The proprietary period for data obtained by ToO proposals will be assessed on a case-by-case basis but will be no more than six months.

2. DDT Exploratory Time. Exploratory Time proposals are normally for requests of small amounts of time, typically a few hours or less, in response to a recent discovery, possibly to facilitate future submission of a larger proposal. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. The possibility that a proposer forgets about or misses an NRAO proposal deadline, or just discovered that he/she was granted time for a particular source on some other telescope, will not constitute sufficient justification for granting observing time by this process. Thus, Exploratory Time proposals must include a clear description of why the proposal could not have been submitted for normal review at a previous NRAO proposal deadline, and why it should not wait for the next proposal deadline. Exploratory Time proposals are evaluated on the basis of scientific merit by Observatory scientific staff. Those staff also asses the technical feasibility of the Exploratory Time proposals. Notification of the disposition of an Exploratory Time proposal normally will be within three weeks of receipt of the proposal; some of these proposals may be put in a queue such that they may or may not be observed. The proprietary period for data obtained by Exploratory Time proposals normally will be six months.

4.3. Helpdesk

Assistance with proposal submission, observation preparation, archive access, and data reduction using both CASA and AIPS is available through the NRAO Helpdesk.  We try our best to get back to you within one working day, although it can take repeated back-and-forth communication to completely resolve your question.  You log on with your my.nrao.edu user name and password.  If you don't have an account you will be given the opportunity to register.

4.4. Observation Preparation

To observe with the VLA on an approved proposal, scheduling blocks must be prepared using the "Observation Preparation Tool," or OPT. The OPT is available at:


Once an observing proposal has been approved for scheduling, about 2-4 weeks before the relevant array configuration in the observing semester, the details of the proposal will appear as a project with program blocks in the OPT. Source and resource catalogs will also be populated as much as possible. The proposers will be notified by email when this has occurred, and when the scheduling queue will be accepting scheduling blocks. Please adhere to the current restrictions posted on https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/restrictions;  things do change and the current restrictions page reflects our latest requirements for accepting scheduling blocks.

After consulting the documentation, any remaining questions or help required in making scheduling blocks can be directed to the NRAO Helpdesk at http://help.nrao.edu using your my.nrao.edu/PST user account.

4.5. Dynamic and Fixed-Date Scheduling

Most projects on the VLA are observed dynamically, based on a combination of the scheduling priority and the expected properties of the array and the weather. Such dynamic scheduling enhances science data quality and the array's ability to discharge time-sensitive science. Some projects requiring coordinated observations with other telescopes will be scheduled on a fixed date.

For further information, see the scheduling officers' home page at http://science.nrao.edu/facilities/evla/sched/schedsoc

4.6. The Observations and Remote Observing

Since most VLA observations are carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Reservations for the VLA site and/or SOC for information on coming to and staying in Socorro.

For those who choose to process their data at home, the data can be retrieved from the VLA online archive by using the NRAO Interactive Services login from the archive web page.  Depending on the size of your data and the speed of your transmission you may choose to retrieve your data online or have us mail you a disk containing your data.

For more information we refer to the VLA archive web page and our Data Shipment Policy.

4.7. Data Access

The online archive contains all VLA data since observing started in 1976, and its contents are available via the Archive Access Tool (AAT) . This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user-specified criteria. Data can be downloaded via standard ftp protocols and via shipment of physical harddisks, subject to the conditions of our disk shipment policy. With the exception of some rapid response and large proposals, VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by any member of the observing team by logging in to the AAT with the my.nrao.edu/PST user account or by making use of the project key (see The Observations and Remote Observing).

VLA data taken after January, 2010, when the transition to the WIDAR correlator took place, are stored in the Science Data Model (SDM) format that is used by both the VLA and ALMA. VLA data are available through the AAT in the following formats:

  • In the native Science Data Model (SDM) format.
  • As a CASA Measurement Set (MS) created from the SDM, which may be averaged in frequency and/or time, to be used with CASA.

Since the start of D configuration in 2013 the archive also contains pipeline calibrated MS.  Though eventually these data will be made accessible through the AAT as well, for the time being they have to be requested through our helpdesk

Some archival data have known problems which, together with the possible fix, are listed at the archive issues page.  Delivery of data in the UVFITS format has been discontinued in June 2013. Instructions to create UVFITS from the SDM or CASA format are given below.

Creating UVFITS or loading data into AIPS using the SDM format

Observers who have installed Obit can convert SDM data directly into UVFITS or load SDM data directly into AIPS.  Note that Obit cannot convert complex correlator modes into a single UVFITS/AIPS file, but plain continuum or simple line data is not a problem.

The latest version of Obit can be found by scrolling to the bottom of the Obit distribution page. Be aware that Obit is a bit harder to install on Macs than on Linux architectures.

To convert SDM data into UVFITS use these steps:

  1. download the SDM format from the archive
  2. using Obit run the task "ASDMList" to obtain a correlator configuration listing for selection
  3. using Obit run the task "BDFIn" to specify the output UVFITS file name

Alternatively use the AIPS path below and write out the file using the AIPS task "FITAB" (or FITTP)

To load SDM data into AIPS use these steps:

  1. download the SDM format from the archive
  2. using AIPS run the task "BDFLIST" to obtain a correlator configuration listing for selection
  3. using AIPS run the task "BDF2AIPS" to specify the output AIPS file name and disk number

Creating UVFITS formatted data using CASA

To produce UVFITS formated data using CASA take the following steps:

  1. download the data from the archive either in SDM or in MS format
  2. Start CASA
  3. Unless you downloaded MS data, convert the SDM formatted data to MS using the CASA task importvla
  4. Apply the requantizer gains as described in the section on 3-bit data observing in the VLA Observing Guide
  5. Convert this MS file to UVFITS format using the CASA task exportuvfits

4.8. Data Processing

The primary data reduction package for the VLA is the CASA (Common Astronomy Software Applications) package, which is also used for ALMA.  See http://casa.nrao.edu for more information on the latest release.

It is possible to reduce VLA data taken in all correlator modes that are offered for general observing using CASA.  More complex correlator set-ups are generally better handled by CASA, and advanced imaging algorithms are available in CASA.  The VLA calibration pipeline uses CASA.

The Astronomical Image Processing System (AIPS) can reduce all observations made with simple correlator set-ups (e.g., identical spw bandwidths and channel widths throughout).  AIPS is currently required for the reduction and analysis of phased-array VLA observations for VLBI and low-band observations (P-band, 230-470 MHz). NRAO will continue to support AIPS for the foreseeable future, at least until VLBI functionality and the full data path for frequencies below 500 MHz have been incorporated into CASA.  See http://www.aips.nrao.edu for more details.

4.9. Travel Support for Visiting the SOC and VLA

Travel support is available for U.S.-based observers to travel to NRAO to observe and/or reduce data.  Please see Travel Support for NRAO Observing Runs and Data Reduction for Non-NRAO Employees for details of eligibility and how to claim reimbursement.

4.10. Student Observing Support Program

In addition to travel support for individual data reduction visits NRAO maintains a program to support research by students, both graduate and undergraduate, at U.S. universities and colleges. Regular and Large proposals submitted for the VLA, VLBA, and GBT, and any combination of these telescopes, are eligible. New applications to the program may be submitted along with new observing proposals at any proposal deadline. Details of this program can be found at https://science.nrao.edu/opportunities/student-programs/studentprograms

4.11. Computing at the SOC

A primary goal of the computing environment at the SOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the SOC for full-time data reduction by visitors. They are mostly two-processor, 8 core , 2.4-GHz PCs with 24GB of memory and 4TB of local disk running Linux.

For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer.

Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/facilities/evla/visitorinfo/visitorregform (see also Reservations for the VLA site and/or SOC). Note that users may request remote access to the visitor machines as well, without actually visiting the SOC.  Remote reservations are subject to availability. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the SOC.

For a more complete description of computing facilities at the SOC, see http://www.aoc.nrao.edu/computing/.

4.12. Reservations for the VLA site and/or SOC

Accommodation for visitors is no longer available at the VLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil.

Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks' notice is preferred, through the online form at http://science.nrao.edu/facilities/vla/visitorinfo/visitorregform. Computing requirements and the level of staff assistance needed must be specified through the online form.

First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator.

4.13. Staying in Socorro

Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at the visitor registration form.

Since Socorro is located just off I-25, it has more than its share of motels.  Most are located on California Street, the main business thoroughfare.  Typical walking distance to NRAO is 20 - 30 minutes.

4.14. Help for Visitors to the VLA and SOC

We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their observing. While in Socorro, each observer will interact with members of the SOC staff in accordance with his/her level of experience and the complexity of the observing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions.

4.15. On-Line Information about the NRAO and the VLA

NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the VLA may be found at http://science.nrao.edu/facilities/evla/.  These pages are regularly updated by the NRAO staff.

5. Publication Guidelines

5.1. Acknowledgement

Any papers using observational material taken with NRAO instruments (Karl G. Jansky Very Large Array or otherwise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF:

The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.

When referring to the VLA in a publication for the first time, please use the full name Karl G. Jansky Very Large Array (VLA).  In subsequent references in that same publication, it is OK to use the shorthand VLA.

5.2. Dissertations

Students whose dissertations include observations made with NRAO instruments are expected to provide copies of, or links to, their thesis for inclusion and maintenance at the NRAO library. These will be cataloged and made available via the NRAO library catalog.

5.3. Preprints

NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO author(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu).

5.4. Reprints

Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers.

5.5. Page Charge Support

The following URL contains complete information on the observatory's policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml.   In summary:

  • When requested, NRAO will pay the larger of the following:
    • 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s).
    • 100% of the page charge share for NRAO staff members.
  • Page charge support is provided for publication of color figures.
  • To receive page charge support, authors must comply with all of the following requirements:
    • Include the NRAO footnote in the text (see Acknowledgement to NRAO).
    • Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254.

6. Documentation

Documentation for VLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see On-Line Information about the NRAO and the VLA). 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.
  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/cont);
    2. SPECTRAL LINE OBSERVING (https://science.nrao.edu/facilities/evla/observing/spectral-line-observing); 
    3. HIGH FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/line);
    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/Doc/Cookbook/casa_cookbook.pdf.
  5. *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.
  6. *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.
  7. *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 http://www.vla.nrao.edu/astro/calib/manual/.
  8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Telescope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983.
  9. *HISTORICAL VLA MEMO SERIES: archive memo series from the early days of the VLA.  See http://www.vla.nrao.edu/memos/.
  10. *RECENT VLA MEMO SERIES: the memo series relating to the expanded capabilities of the VLA.   See http://www.aoc.nrao.edu/evla/memolist.shtml.
  11. *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.
  12. 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 the best source of information on much of the VLA. Copies of this are found at the VLA and in the SOC, but no new copies are available. Much of this document is available for downloading through the NRAO's website. It does not include any information about the hardware and software specific to the expanded Karl G. Jansky VLA.

7. Key Personnel

Please direct queries to the NRAO Helpdesk, at http://science.nrao.edu/observing/helpdesk.shtml, and they will be able to answer your query or direct it to the appropriate person for followup.  

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 the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835.

8. Acknowledgements

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. Please contact the editor of the present document (gvanmoor at nrao dot edu, dshepher at nrao dot edu) with questions on the material, or suggestions that would enhance the clarity of this guide.