VLA > VLA Observational Status Summary 2015A

# VLA Observational Status Summary 2015A

VLA capabilities February 2015 - September 2015

# 1.1. Purpose of Document, Older Versions of the OSS

This Observational Status Summary (OSS) summarizes the instrumental status of the Karl G. Jansky Very Large Array (VLA) for the B, BnA, and A configurations starting 6 February 2015, and should be used when preparing proposals for the 1 August 2014 deadline.  For capabilities prior to that date, we refer to our overview of all OSS versions available online.

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

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

# 1.2. 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 2014/2015 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 the Guide to Proposing for the VLA for information on how to submit an observing proposal.

Table 3: VLA Configuration Schedule for 2014/2015

27 Jun - 17 Sep, 2014 19 Sep - 6 Oct, 201417 Oct - 5 Jan, 2015 9 Jan - 26 Jan, 20156 Feb - 11 May, 2015 15 May - 1 Jun, 201512 Jun - 21 Sep, 2015
Configuration: D DnC C CnB B BnA A

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

All VLA antennas are outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers cover the frequency ranges of 1-2 GHz, 2-4 GHz, 4-8 GHz, 8-12 GHz, 12-18 GHz, 18-26.5 GHz, 26.5-40 GHz, and 40-50 GHz and the bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively.  In addition, all antennas of the VLA now have even lower frequency receivers, covering 230-470 MHz (P-band).  And a small number are being outfitted with even lower frequency receivers, covering 54-86 MHz (4-band).

The VLA correlator is both powerful and flexible.  Details of the correlator configurations being offered for VLA science in 2015 are described in the WIDAR Section of this document.   It is important to realize that the VLA correlator is fundamentally a spectral line correlator and that even 'continuum' observations are done in a wide-band mode with many channels.

# 1.3. The Expanded Very Large Array (EVLA) Project

The Expanded VLA (EVLA) project modernized the VLA electronics (built in the 1970s and 1980s) in order to improve several key observational parameters by an order of magnitude or more.  Some of the details of the EVLA Project may be found at http://www.aoc.nrao.edu/evla/.   The EVLA project was funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico.   Total funding was approximately $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 capable of processing the large bandwidth with a minimum of 16,384 spectral channels per baseline.  All goals were met.   A comparison of some of the EVLA performance parameters with those of the original VLA is provided in Table 1.

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

VLA to EVLA transition

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

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

# 2. Proposing for the VLA for the Next Semester

### The 2015A Call for Proposals

The 2015A Call for Proposals details the General Observing (GO) capabilities being offered for the Karl G. Jansky Very Large Array (VLA).

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

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

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

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

#### Summary of Capabilities

As described in the 2015A Call for Proposals, the VLA offers continuous frequency coverage from 1-50 GHz in the following observing bands: 1-2 GHz (L-band); 2-4 GHz (S-band); 4-8 GHz (C); 8-12 GHz (X); 12-18 GHz (Ku); 18-26.5 GHz (K); 26.5-40 GHz (Ka); and 40-50 GHz (Q).  Both single pointing and mosaics with discrete, multiple, field centers will be supported.  Data rates of up to 60 MB/s (216 GB/hour) will be available to all users, combined with correlator integration time limits per band and per configuration, as described in the Time Resolution and Data Rates.  Data rates in excess of 25 MB/s require additional justification.  Limitations on frequency settings and tuning ranges are described in the Frequency Bands and Tunability Section.

The GO capabilities being offered are:

 Capability Description 8-bit samplers Standard default set-ups for: 2 GHz bandwidth continuum observations at S/C/X/Ku/K/Ka/Q bands (16 x 128 MHz sub-bands) 1 GHz bandwidth continuum observations  at L-band (16 x 64 MHz sub-bands) 256 MHz bandwidth continuum observations at P band (16 x 16 MHz sub-bands) Flexible set-ups for spectroscopy, using two independently tunable 1 GHz basebands, each of which can be split into up to 16 flexibly tunable sub-bands Single, dual, and full polarization products Number of channels summed over all polarization products up to 16,384 (no recirculation) or up to 65,536 (with recirculation) 3-bit samplers Standard default set-ups for: 8 GHz bandwidth continuum observations at K/Ka/Q bands 6 GHz bandwidth at Ku band 4 GHz bandwidth at C/X bands Flexible set-ups for spectroscopy, using four independently tunable 2 GHz basebands, each of which can be split into up to 16 flexibly tunable sub-bands Single, dual, and full polarization products Number of channels summed over all polarization products up to 16,384 (no recirculation) or up to 65,536 (with recirculation) Mixed 3-bit and 8-bit samplers Allows more flexibility for simultaneous continuum and high-resolution spectral line observing Sub-arrays Up to 3 independent sub-arrays using standard 8-bit continuum set-ups Phased-array for VLBI VLA Phased Array ("Y27") for VLBI (see the VLBA call for proposals and the VLBA Observational Status Summary for a detailed description of the VLA Phased Array capabilities being offered)

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

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

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

#### Steps to Completion

1. Read the NRAO 2015A Call for Proposals description for summaries of the capabilities being offered for GO and SRO.
2. If you are proposing for spectral line observations, run the GO Setup Tool (GOST, a java application) to define your correlator set-up and save a snapshot of the tool GUI to disk.  Note:
• Write a scientific justification for your proposal.  Note that technical information should be included in the Technical Justification section and does not need to be included with the scientific justification.  Save your scientific justification as a PDF file.
• Log into the NRAO Interactive Services page (my.nrao.edu) and click on the "Proposals" tab in the top left to create a new proposal for the VLA.  Once you are in the tool, extensive help is available in the tool by clicking on the "Help" button in the top right of the tool interface (see also the links to the help documentation below).  In the PST:
• Fill in the relevant fields for each section of your proposal;
• If you are proposing for continuum observations, select Observing Type = "Continuum" in the General section of your proposal.  Then add a Continuum Resource, selecting the appropriate band you want to observe in.  Default optimum continuum setups for each band are defined in the PST.
• If you are proposing for spectral line observations, Select Observing Type = "Spectroscopy" in the General section of your proposal.  Then add a Spectroscopic Resource and attach the GOST snapshot (that you saved previously) to this resource.  The GOST snapshot describes everything needed to specify the lines you want to observe and what correlator resources are needed for the observation.
• Technical justifications are now included as a separate section in the proposal.  Click on the Technical Justification page in the proposal and fill in the appropriate boxes. In order to determine sensitivities, you will need to use the VLA Exposure Calculator Tool (ECT) to understand what sensitivity you will get for a given frequency, bandwidth, and integration time.  Some guidelines for running the ECT can be found here. If you have trouble running the ECT please contact the NRAO helpdesk;
• Note that the PST allows observers to specify multiple resource types (e.g., you can have one proposal that specifies general and/or RSRO correlator resources).  If any resource is RSRO, it will be a RSRO proposal.
• When your proposal is complete, validate it to make sure there are no obvious omissions or mistakes.
• When you are satisfied, submit your proposal in the PST.
•

### Resident Shared Risk Observing (RSRO)

#### Summary of Capabilities

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

• correlator dump times shorter than 50 msec, including integration times as short as 5 msec for transient detection
• pulsar observations
• data rates above 60 MB/s
• recirculation beyond a factor of 64 in the correlator
• P-band system polarimetry and spectroscopy
• more than 3 sub-arrays or sub-arrays with the 3-bit system, and
• complex phased array observations (e.g., pulsar and complex VLBI observing modes).

#### Steps to Completion

1. Read the 2015A NRAO Call for Proposals for a summary of the capabilities being offered for GO and SRO. If you want more than what is offered for GO or SRO then you are requesting a RSRO capability (one that is not well-tested or may even need additional development).  If you propose for a RSRO capability, you (or an experienced person on your team) must be able to participate in the RSRO program by coming to Socorro to help with the development and testing process.
2. Write a scientific justification for your proposal.  Note that technical information should be included in the Technical Justification section and does not need to be included with the scientific justification.  Save your scientific justification as a PDF file.
3. Log into the NRAO Interactive Services page (my.nrao.edu) and click on the "Proposals" tab in the top left to create a new proposal for the VLA.  Once you are in the tool, extensive help is available in the tool by clicking on the "Help" button in the top right of the tool interface (see also the links to the help documentation below).  In the PST:
• Fill in the relevant fields for each section of your proposal;
• Even as a RSRO proposal, you will need to create a "Resource" in the PST and select the "WIDAR RSRO" Back End.  This will give you a text field in which you can type a description of your set-up;
• Technical justifications are now included as a separate section in the proposal.  Click on the Technical Justification page in the proposal and fill in the appropriate boxes. Describe what you want to do with the correlator and why this is in the RSRO Category in this section. Contact NRAO staff if you have questions about exactly what is feasible. The last box at the bottom of the Technical Justification page should be used to describe your RSRO effort. Identify who in your team can come to Socorro to help commission this capability and how their background and expertise can be applied to this development.  Working with NRAO staff, estimate the level of effort that is likely to be needed for this development and specify how long a member of your team can come to NRAO.
• Note that the PST allows observers to specify multiple resource types (e.g., you can have one proposal that specifies general, shared-risk and RSRO correlator resources).  If any resource is RSRO, it will be a RSRO proposal.
• When your proposal is complete, validate it to make sure there are no obvious omissions or mistakes.
• When you are satisfied, submit your proposal in the PST.
• #### Considerations

A RSRO proposal should contain:

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

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

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

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

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

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

RSRO participation without a science proposal:

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

### Documentation Summary

https://science.nrao.edu/enews/7.7/ - The July 7, 2014 NRAO call for proposals for the August 1, 2014 deadline (2015A).

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

https://my.nrao.edu/nrao-2.0/PSTMANUAL/PSTMANUAL.html - The NRAO Proposal Submission Tool (PST) online manual. Also available in PDF format.

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

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

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

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

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

# 3.1. Introduction

This section contains details of the VLA's resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in the 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.
Footnotes:
1. Bmax is the maximum antenna separation, Bmin is the minimum antenna separation, θHPBW is the synthesized beam width (FWHM), and θLAS is the largest scale structure "visible" to the array.
2. The listed resolutions are appropriate for sources with declinations between −15 and 75 degrees. For sources outside this range, the extended north arm hybrid configurations (DnC, CnB, BnA) should be used, and will provide resolutions similar to the smaller configuration of the hybrid, except for declinations south of −30. No double-extended north arm hybrid configuration (e.g., DnB, or CnA) is provided.
3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for θHPBW. The values for snapshots are about 1.3 times the listed values.
4. The largest angular scale structure is that which can be imaged reasonably well in full synthesis observations. For single snapshot observations the quoted numbers should be divided by two.
5. For the C configuration an antenna from the middle of the north arm is moved to the central pad "N1". This results in improved imaging for extended objects, but may slightly degrade snapshot performance. Note that although the minimum spacing is the same as in D configuration, the surface brightness sensitivity and image fidelity to extended structure is considerably inferior to that of the D configuration.

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

# 3.3. Sensitivity

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

where:

- SEFD is the "system equivalent flux density" (Jy), defined as the flux density of a radio source that doubles the system temperature. Lower values of the SEFD indicate more sensitive performance. For the VLA's 25-meter paraboloids, the SEFD is given by the equation SEFD = 5.62TsysA, where Tsys is the total system temperature (receiver plus antenna plus sky), and ηA is the antenna aperture efficiency in the given band.
- ηc is the correlator efficiency (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 1 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 = 0.78 to 0.83.
In general, the zenith atmospheric opacity to microwave radiation is very low - typically less than 0.01 at L, C and X-bands, 0.05 to 0.2 at K-band, and 0.05 to 0.1 at the lower half of Q-band, rising to 0.3 by 49 GHz. The opacity at K-band displays strong variations with time of day and season, primarily due to the 22 GHz water vapor line. Conditions are best at night, and in the winter. Q-band opacity, dominated by atmospheric O2, is considerably less variable.
Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Significant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz.
Tipping scans can be used for deriving the zenith opacity during an observation. In general, tipping scans should only be needed if the calibrator used to set the flux density scale is observed at a significantly different elevation than the range of elevations over which the phase calibrator and target source are observed. However, the antenna tip capability is currently unavailable for the upgraded VLA -- it is hoped that this will again be available in the next year.

When the flux density calibrator observations are within the elevation range spanned by the science observing, elevation dependent effects (including both atmospheric opacity and antenna gain dependencies) can be accounted for by fitting an elevation-dependent gain term. See the following item.
• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target.

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

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

To aid VLA proposers there is an on-line Exposure Calculator tool 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:

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):

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

# 3.4. VLA Frequency Bands and Tunability

### Bands

For observations taken with the 8-bit samplers,  each receiver can tune to two different frequencies, each 1.024 GHz wide,  within the same frequency band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies, except for the low-band receiver (50 -- 500 MHz), which provides linear polarization.  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.   In 8-bit mode, a maximum of 1.024 GHz can be correlated for each IF pair (see the WIDAR Section), for a total maximum bandwidth of  2.048 GHz.  To distinguish this 8-bit system from the 3-bit system (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.

Notes:
1.  Listed here are the nominal band edges.  For all bands, the receivers can be tuned to frequencies outside this range, but at the cost of diminished performance.  Contact VLA staff for further information.
2.  The 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.

### Tuning Restrictions

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, with A0C0 replaced by A1C1 and A2C2, and B0D0 replaced by B1D1 and B2D2.

# 3.5. VLA Samplers

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

The 8-bit Set consists of four 8-bit samplers running at 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.  Some restrictions apply, depending on band, as described in the section on Frequency Bands and Tunability

The 3-bit Set consists of eight 3-bit samplers running at 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.  The tuning restrictions are described in the section on Frequency Bands and Tunability.

Which set to use?

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

Major Characteristics of each Set

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

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

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

Setting up the 8-bit or 3-bit Samplers

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

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

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

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

Other issues

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

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

# 3.6.1. 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

# 3.6.2. 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.

# 3.6.3. 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.

# 3.6.4. 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
Bands
Default
integration time
D, C
L S C
5 seconds
D, C
X Ku K Ka Q
3 seconds
B all 3 seconds
A all 2 seconds

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

Minimum Integration Times and Maximum Data Rates
(2014A)
Proposal type

Minimum integration time

Maximum data rate
Standard 50 msec
25 MB/s ( ~88 GB/hr)
Shared risk (SRO) 50 msec 60 MB/s (~210 GB/hr)

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

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

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

 Data rate = 45 MB/sec × (Nchpol/16384) x Nant × (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 ~2 seconds for standard observations, and 0.8 seconds for shared risk observations.

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

Higher time resolutions and data rates are possible in principle but will be considered only through the

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

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

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

# 3.9. 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

# 3.10. Positional Accuracy & Astrometry

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

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

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

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

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

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

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

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

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

The dominant widefield effects are differential aberration and EOP errors. Data exported from the EVLA archive as a CASA measurement set will have reasonably accurate EOP attached by the CASA FILLER, based on BIPM values pertaining to the observation date, but determined after the fact, on the date of export. These EOP differ slightly from those used at correlation (observation) time, when only predictions are available. The difference is usually negligible.  However, CASA (program FIXVIS) does not, as of June 2013, correct for the differential aberration.    On the other hand, AIPS corrects differential aberration (program UVFIX), but, because there is no EOP data in the archived ASDM, the UVFITS file imported by BDF2AIPS has zeros in the antenna table where the EOP should be. One can edit in the EOP and proceed with UVFIX, but it is better that the  ASDM  stores EOP in the first place.  This feature is expected by late 2013.  The capabilities of  both CASA and AIPS should be equivalent in this regard by 2014.

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

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

# 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 non-Deconvolved Sources

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

# 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, and is also available within the Observation Preparation Tool.

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

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

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

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

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.

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

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

• 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 - a couple of times per hour is generally sufficient.
• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale.  The 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.  A detailed description of calibrator codes is available in the Key to the calibrator list.
• In clear and calm conditions, most notably in the summer, phase stability often deteriorates dramatically after about 10AM, due to small-scale convective cells set up by solar heating.  Observers should consider a more rapid calibration cycle for observations between this time and a couple hours after sundown.
• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See Rapid Phase Calibration and the Atmospheric Phase Interferometer (API).
• Use 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 for very small source-calibrator separations. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the old VLA this was known as "fast-switching." For the upgraded VLA it is just a loop of source-calibrator scans with short scan length. This technique "stops" tropospheric phase variations at an "effective" baseline length of ∼vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest VLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Note, however, that for the compact D-configuration, and a typical wind velocity, this "effective" baseline length is the same as, or larger than, the longest baseline in the array, and it is not worth the increased overhead of short cycle times. Under these circumstances it is sufficient to calibrate every 5-10 minutes to track the instrumental changes. The fast switching technique will also not work in bad weather (such as rain showers, or when there are well-developed convection cells - most notably, thunderstorms). It is also important to specify correctly the required tropospheric phase stability as measured by the Atmospheric Phase Interferometer at observe time (see below).

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

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

Characteristic seasonal averages are represented below:

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.

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

33

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

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

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

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

# 3.15. VLBI Observations

The VLA can now participate in VLBI observations with the VLBA.  Currently, this is only allowed in phased array mode (single dish is only available through the VLBA Resident Shared Risk Observing program) with restricted WIDAR correlator configurations.   For more details see the VLBI at the VLA documentation.  In phased array mode the program TelCal derives the antenna-based delay and phase corrections needed for antenna phasing in real time.  This correction is applied to the antenna signals before they are summed, requantized to 2-bits, and recorded in VDIF format on the Mark5C disk at the VLA site.  The disk(s) are then transported to Socorro, NM and correlated on the DiFX correlator with other VLBI stations which participated in the observation.  Standard VLA data, i.e., correlations between VLA antennas,  are also archived in the NRAO science data archive.

# 3.16. 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://cornish.leeds.ac.uk/public/index.php)" (B configuration), or the NRAO VLA Sky Survey (NVSS, http://www.cv.nrao.edu/nvss/) (D configuration, all-sky) surveys.

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.

# 3.18. Combining Configurations and Mosaicking

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

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

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

Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicking.  The VLA has no limit on the number of pointings for each mosaic. Typically hexagonal, rectangular, or individual pointing patterns are used and the overlap regions will result in an improved rms over each individual pointing. Given the many, potentially short observations, it is important to obey the data rate limits outlined in the Time Resolution and Data Rates Section.

As of Semester 2015A on-the-fly mosaics, i.e. dumping the data while moving the telescopes across the source, is available as Shared Risk Observing.

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 mosaicking with the VLA are given in the VLA Observing Guide.

# 3.19. Pulsar Observing

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

# 4.1. Introduction

The correlator configurations offered for general observing in the 2014B 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 2014B semester.

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.

# 4.2. Wideband Observing

The wideband observing setups provide the widest possible bandwidth for a given observing band, with channel spacing depending on the number of polarization products.

Wideband & Subarray Correlator Options(all but L band)

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

8-bit wideband setups are available for all observing bands, providing a total of 2 GHz of bandwidth per polarization (1 GHz per polarization at L-band).  3-bit setups are available for all bands above S-band, providing total bandwidths per polarization of 4 GHz (C/X bands), 6 GHz (Ku band), or 8 GHz (K/Ka/Q bands). In all cases but L band each of the subbands is 128 MHz wide (at L band the default is 64 MHz/subband, yielding channels twice as narrow as those listed in the table above). Recent tests suggest use of the 3-bit samplers yields higher continuum sensitivity; see the section on VLA Samplers for details.

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.

# 4.3. Spectral Line Observing

### Basebands and Subbands

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

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

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

### Subband tuning restrictions

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

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

where:

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

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

The figure below illustrates these restrictions:

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

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

### Subband bandwidths & the digital filter response

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

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

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

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

### Spectral channels and polarization products

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

1. For the 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, 3, 4, 5, ..., 64 is the number of Baseline Board Pairs (BlBPs) assigned to that subband.
2. The sum over all subbands of nBlBP,i must be less than or equal to 64, the number of Baseline Board pairs in the correlator. Equivalently, the sum over all subbands of spectral channels times polarization products is limited to 64x256= 16,384 (without recirculation).

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

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

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

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

Channel spacing:
128 MHz 38400/νGHz km/s 2000/nBlBP kHz 600/nBlBPGHz km/s 1000/nBlBP kHz 300/nBlBPGHz km/s 500/nBlBP kHz 150/nBlBPGHz km/s
64 19200 1000 / nBlBP 300 / nBlBP 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP
32 9600 500 / nBlBP 150 / nBlBP 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP
16 4800 250 / nBlBP 75 / nBlBP 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP
8 2400 125 / nBlBP 37.5 / nBlBP 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP
4 1200 62.5 / nBlBP 18.75 / nBlBP 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 /n BlBP
2 600 31.25 / nBlBP 9.375 / nBlBP 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP
1 300 15.625/nBlBP 4.687 / nBlBP 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP
0.5 150 7.8125 / nBlBP 2.344 / nBlBP 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP
0.25 75 3.906 / nBlBP 1.172 / nBlBP 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP
0.125 37.5 1.953 / nBlBP 0.586 / nBlBP 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP
0.0625
18.75 0.977 / nBlBP 0.293 / nBlBP 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP
0.0325 9.375 0.488 / nBlBP 0.146 / nBlBP 0.244 / nBlBP 0.073 / nBlBP 0.122 / nBlBP 0.037 / nBlBP
Subband bandwidth and spectral resolution options. Note that the table entries refer to the spacing between spectral channels -- that spacing is before any frequency smoothing, so these channels are not independent.
• nBlBP is the number of Baseline Board Pairs assigned to the subband.
• Each subband may have a different number of spectral channels and polarization products, and each may be tuned independently.
• There can be at most 16 subbands per baseband, and nBlBP must be an integer: 1, 2, 3, 4, 5, ..., 64.
• The sum of nBlBP over all subbands must be less than or equal to 64.
• Use of more than one BlBP for a subband may further restrict the number of subbands available in one or more of the basebands; see text for details.

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

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

### Recirculation

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

At some time in the future an alternative method of recirculation, using extra the CPU cycles freed up by increasing the integration time, will be made available for General Observing as well. Recirculation by limiting of the subband bandwidth to increase the number of channels (in factors of two), was used in the pre-expansion VLA correlator.

#### Recirculation versus Baseline Board stacking

When faced with the choice between Recirculation and Baseline Board stacking (see below) to increase the number of channels in a subband for subbands narrower than 128 MHz we recommend the former, which is supported in observatory software (GOST, OPT).

The current implementation of Recirculation is that for each halving of the subband bandwidth the number of channels in the subband may be doubled without having to trade off the use of other subbands. Because recirculation is achieved by limiting the subband bandwidth, it is not supported for 128 MHz subbands, whereas for 64 MHz subbands only a factor 2 recirculation is supported, etc. The maximum recirculation factor for a subband is 128/(subband bandwidth in MHz), and of course also subject to other configuration restrictions such as data rate.

The juggling between the requested number of channels, subband bandwidth and the available number of Baseline Board pairs is dependent on the science goals and not easily formulated in a standard answer. However, if subbands of less than 128 MHz are used, Recirculation becomes an option for setups that can also be achieved with Baseline Board stacking. In such cases we suggest to use Recirculation where possible, and within the General Observing or Shared Risk requirements.  This frees up unused Baseline Board pairs for other use; alternatively, one becomes less dependent on all Baseline Board pairs being in working order.

Recirculation with factors 8 to 64 is designated Shared Risk, and recirculation with factors over 64 Resident Shared Risk. The latter choice may have severe implications for the sensitivity as visibility integration time is used as trade-off. Ask the NRAO User Helpdesk for more details.

### Baseline Board stacking

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

#### Baseline Board Stacking and Correlator Use

First let us consider how the correlator hardware is organized. The cross-multiplications in the WIDAR correlator are spread across 64 Baseline Board (BlB) pairs, arranged into 4 quadrants of 16 BlB pairs each. Each baseband is connected directly to one of those quadrants. In the simplest mode, each of the 16 BlB pairs of a quadrant handles the correlations for one of the 16 subbands of the corresponding baseband. Four basebands and four quadrants are required, in order to handle the full 8 GHz of bandwidth per polarization provided by the 3-bit (wideband) samplers: that 8 GHz is split into four basebands of 2 GHz each, with each baseband fed into a different BlB quadrant. Each BlB pair in that quadrant handles a subband of maximum bandwidth 128 MHz, so 16 BlB pairs handles 16 subbands for a total of 16x128 MHz= 2048 MHz.

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

When using the 8-bit samplers, the total bandwidth is only 2 GHz per polarization, split into two basebands of 1 GHz each. The simplest continuum setup uses only two quadrants, since there are only two basebands; and only 8 subbands are required to span the 8x128 MHz= 1024 MHz of each baseband. Three-quarters of the correlator BlB hardware remain unused.

The spectral line mode allows access to these `extra' correlator resources through Baseline Board stacking: using multiple BlB pairs to process the same subband and produce more cross-correlations for that subband. This is done using crossbar switches which make the data for a single subband available to several BlB pairs. Those BlB pairs can then be used to produce more spectral channels for that subband, with n BlB pairs producing 256*n cross-correlations per baseline. The limit on the total number of cross-correlations (16384) stems from the total number of BlB pairs (64): 64x256= 16384.

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

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

Routing option #1 means that one could use all the BlB pairs within a quadrant to correlate a single subband, yielding 16x256= 4096 cross-correlations for that subband:

Routing option #2 means that one could use the BlB pairs in all 4 quadrants to correlate a single subband. One simple case would use 4 BlB pairs to correlate each of the 16 subbands in a single baseband, yielding 4x256= 1024 cross-correlations for each of those subbands. Note that in this case, no BlB pairs are left to correlate any data from the second baseband.

Using routing option #2 does come with a subtle cost: assigning a BlB pair in quadrant X to correlate a subband corresponding to quadrant Y, removes that BlB pair from use in the baseband corresponding to quadrant X...and hence also removes the corresponding subband in that baseband. So getting more channels for a subband in one baseband, may prevent the use of a subband in a different baseband. To take a simple example, consider an experiment where one wishes to observe a single line in dual polarization with 512 channels (requiring 4 BlB pairs), plus as much continuum bandwidth as possible. Naively one would say there are 16 subbands in each baseband; one is used for the spectral line, so that leaves 16+15= 31 subbands, and with the widest subband bandwidth (128 MHz) the total available continuum should be 31x128 MHz= 3968 MHz per polarization. Actually however there are only 15+15 subbands available, or 30x128 MHz= 3840 MHz per polarization, because the spectral line subband has "eaten" one BlB pair corresponding to the other baseband:

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

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

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

The above examples all use BlB pair stacking in powers of 2, but this is not required.  To give some idea of more complex possibilities, the following tables give two examples of other possible configurations. The 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.

Complex 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:

Complex 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:

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

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

3-bit Complex Configuration Example #1
BasebandSubbandPol'n products Spectral channelsnBlBP
A1/C1 sb0-8 RR, LL, RL, LR 9 x 64
9 x 1
A1/C1 sb9 RR, LL 1 x 1152 1 x 9
A1/C1 sb10 RR 1 x 1792
1 x 7
A1/C1 sb11 RR, LL 1 x 384
1 x 3
A2/C2 sb0-11 RR, LL, RL, LR 12 x 64 12 x 1
A2/C2 sb12 LL 1 x 768 1 x 3
B1/D1 sb0-3 RR, LL, RL, LR 4 x 64
4 x 1
B1/D1 sb4 RR, LL, RL, LR 1 x 320
1 x 5
B2/D2 sb0-6 RR, LL, RL, LR 7 x 64 7 x 1
B2/D2 sb7 RR, LL 1 x 640 1 x 5
SRCT display:

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.

# 5. Documentation

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

1. PROCEEDINGS FROM THE 1988 SYNTHESIS IMAGING WORKSHOP: Synthesis theory, technical information and observing strategies can be found in: "Synthesis Imaging in Radio Astronomy." This collection of lectures given in Socorro in June 1988 has been published by the Astronomical Society of the Pacific as Volume 6 of their Conference Series.
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/vla/docs/manuals/obsguide/modes/line);
3. HIGH FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/hifreq);
4. LOW FREQUENCY OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/lofreq);
5. VERY LOW FREQUENCY OBSERVING (< 500 MHz) (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/vlofreq);
6. POLARIMETRY (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/pol);
7. MOSAIC OBSERVING (https://science.nrao.edu/facilities/vla/docs/manuals/obsguide/modes/mosaicking);
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 https://science.nrao.edu/facilities/vla/docs/manuals/cal.
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 a useful source of information on much of the VLA. There are a few hard copies at the VLA and in the DSOC.  Much of this document is now available for downloading.  Note: it does not include any information about the hardware and software specific to the expanded Karl G. Jansky VLA.

# 6. Key Personnel

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

If you wish to contact an NRAO staff member directly, you may e-mail them by addressing your message to "first initial last name"@nrao.edu. Thus, you may contact Joanne Astronomer at: "jastrono@nrao.edu". The name is truncated to eight characters.   For questions about telescope time allocation, please e-mail "schedsoc@nrao.edu".

The listed four-digit numbers are sufficient for calls made from within the SOC.   If you are calling from outside the SOC, dial 575-835 before dialing the four-digit number. Please also refer to our searchable NRAO Directory

# 7. Editor's Notes

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