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| = Introduction = | | = EVLA Observational Status = |
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| == Purpose of Document ==
| | * [[Observational Status Summary]]: Describes expected OSRO capabilities for the next D→A configuration cycle, September 2011 through January 2013, for use in preparing proposals for the August 1, 2011, proposal deadline, and following (updated 29 Jun 2011). |
| | | * [[Observational Status Summary - Current]]: Describes current OSRO capabilities through the end of the A configuration, September 2011, for use by observers and rapid response science proposers (updated 15 May 2010). |
| This document summarizes the current instrumental status of the Expanded Very Large Array (EVLA). It is intended as a ready reference for those contemplating use of the EVLA for their astronomical research. The information is in summary form – those requiring greater detail should consult the EVLA’s staff members, listed in Section 8, or refer to the manuals and documentation listed in Section 7. Most of the information contained here, and much more, is available on the Web, and can be accessed through the EVLA home page information for astronomers, at http://science.nrao.edu/evla/, and the VLA information for astronomers, http://www.vla.nrao.edu/astro/. These pages will shortly be combined into a single page for the EVLA, and the links in this document will be updated accordingly. A companion document for the VLBA is also available from http://science.nrao.edu/vlba/.
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| The EVLA is a large and complex modern instrument. It cannot be treated as a “black box,” and some familiarity with the principles and practices of its operation is necessary before efficient use can be made of it. Although the NRAO strives to make using the EVLA 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 especially encouraged to enlist the assistance of an experienced colleague or NRAO staff member for advice on, or direct participation in, an observing pro- gram. Refer to Section 5.14 for details. The EVLA will be an extremely flexible instrument, and we are always interested in imaginative and innovative ways of using it.
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| == What is the Expanded Very Large Array? ==
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| The EVLA is the product of a program to modernize the electronics of the Very Large Array (VLA) in order to improve several key observational parameters by an order of magnitude or more. Some of the details of the EVLA Project may be found on the web, at http://www.aoc.nrao.edu/evla/. The EVLA is funded jointly by the US National Science Foundation (NSF), the Canadian National Research Council, and the CONACyT funding agency in Mexico. Total funding is 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 will be completed on time and on budget in 2012, 11 years after it began. Its key observational goals are (1) complete frequency coverage from 1 to 50 GHz; (2) continuum sensitivity improvement by up to an order of magnitude (nearly two orders of magnitude in speed) by increasing the bandwidth from the VLA’s 100 MHz per polarization to 8 GHz per polarization; and (3) implementation of a new correlator that can process the large bandwidth with a minimum of 16,384 spectral channels per baseline. A comparison of some of the EVLA performance parameters with those of the VLA is provided in Table 1. The remaining major milestones for the EVLA are shown in Table 2.
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| == VLA to EVLA Transition ==
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| The year 2010 will be extremely exciting for the EVLA. The correlator that has been the heart of the VLA for three decades was decommissioned on 11 January, 2010. The VLA will be shut down to outside users until early March 2010, during which time hardware willbe transferred from the old correlator to the EVLA correlator and observing modes will be commissioned in preparation for EVLA early science. At the same time the direction of the configuration cycles will also change, from A→B→C→D→A to D→C→B→A→D, in order to facilitate the EVLA correlator commissioning and to limit initial EVLA data rates. When the telescope returns to general use it will be the EVLA.
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| {| border="1"
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| |+ '''Table 1: Overall EVLA Performance Goals'''
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| !Parameter
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| !VLA
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| !EVLA
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| !Factor
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| |-
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| | Continuum Sensitivity (1-σ, 9 hr) || 10 μJy || 1 μJy || 10
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| |-
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| | Maximum BW in each polarization || 0.1 GHz || 8 GHz || 80
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| |-
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| | Number of frequency channels at max. BW || 16 || 16,384 || 1024
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| |-
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| | Maximum number of freq. channels || 512 || 4,194,304 || 8192
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| |-
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| | Coarsest frequency resolution || 50 MHz || 2 MHz || 25
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| |-
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| | Finest frequency resolution || 381 Hz || 0.12 Hz || 3180
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| |-
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| | Number of full-polarization sub-correlators || 2 || 64 || 32
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| |-
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| | Log (Frequency Coverage over 1–50 GHz) || 22% || 100% || 5
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| |-
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| |}
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| Note: The "Factor" gives the factor by which the EVLA parameter will be an improvement over the equivalent VLA parameter.
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| {| border="1"
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| |+ '''Table 2: EVLA Major Milestones'''
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| !Milestone
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| !Target Date
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| | Installation of EVLA correlator subset for early science || 2010 Q1
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| | Shared Risk Observing begins || 2010 Q1
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| | Full EVLA correlator installation || 2010 Q2
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| |-
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| | Last antenna retrofitted || 2010 Q2
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| |-
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| | Last receiver installed || 2012 Q3
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| |-
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| |}
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| = An Overview of the EVLA =
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| The EVLA 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 (DSOC) in Socorro, New Mexico.
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| The basic data produced by the EVLA 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) will allow operation of the array as a single element through coherent summation of the individual antenna signals. This mode will most commonly be used for VLBI observing and for observations of rapidly varying objects, such as pulsars. However, it will not be available initially.
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| The EVLA can vary its resolution over a range exceeding a factor of ∼ 50 through movement of its component antennas. There are four basic arrangements, called configura- tions, 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 de- ployed 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.
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| Traditionally, the VLA completed one cycle through all four configurations in ap- proximately a 16 month period. However, this period will likely change in early 2010 to accommodate commissioning of the EVLA correlator and the onset of EVLA early science. The present best estimate for the EVLA configuration schedule in 2010 and 2011 is presented in Table 3, but prospective users should consult the web page http://science.nrao.edu/evla/proposing/configpropdeadlines.shtml or re- cent NRAO and AAS newsletters for up-to-date schedules and associated proposal deadlines. Refer to Section 5.1 for information on how to submit an observing proposal.
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| {| border="1"
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| |+ '''Table 3: Predicted EVLA Configuration Schedule for 2010-2011'''
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| ! Year
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| ! Feb-May
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| ! Jun-Sep
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| ! Oct-Jan
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| | 2010 || D || C || B
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| |-
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| | 2011 || A || D || C
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| |-
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| |}
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| Observing projects on the EVLA 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 EVLA 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 Section 4.15 for details.
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| All EVLA antennas will eventually be outfitted with eight receivers providing continuous frequency coverage from 1 to 50 GHz. These receivers will cover 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. These bands are commonly referred to as L, S, C, X, Ku, K, Ka, and Q bands, respectively. See Section 3.2 for more details about the availability of new bands.
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| The VLA’s original P-band (300 – 340 MHz) receivers are incompatible with the EVLA’s wideband electronics, so there is at present no P-band observing capability. The NRAO, in cooperation with NRL, is now developing a wideband receiver system which will provide improved P-band performance. Tests of this new system will be carried out during 2010, but there is not yet an implementation date. This new receiver system will also replace the existing 74-MHz (4-band) receivers. It is unlikely that the 74-MHz capability will be available in 2010.
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| The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations to be offered for EVLA early science are described in Section 4.13. It is important to realise that the EVLA correlator is fundamentally a spectral line correlator. The days of separate “continuum” and “spectral line” modes of the VLA correlator are over, and all observations with the EVLA will be “spectral line.” This has implications for how observations are set up, and users who may be used to continuum observing with the VLA are strongly advised to consult Section 4.13.
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| = EVLA Early Science =
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| EVLA early science will be provided by two programs for outside users and one for EVLA commissioning staff. All early science programs will be peer-reviewed. In keeping with a primary construction project goal, the EVLA will continue to be used for science throughout the commissioning of the telescope into full operations in 2013. Observing during this period will thus involve an element of risk associated with the large stepwise increases in throughput bandwidth that will be offered to the community at the start of each new array configuration cycle in 2010, 2011, and 2012.
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| The Open Shared Risk Observing (OSRO) program will provide early science capabilities to the general user community. These capabilities will initially provide a maximum 256 MHz bandwidth that will increase to 2 GHz in mid-2011 and to 8 GHz in 2012. The Resident Shared Risk Observing (RSRO) program will provide these capabilities, and other more powerful ones, much sooner to users who can reside in Socorro and help with the EVLA commissioning efforts. These same enhanced capabilities will also be made available to EVLA commissioning staff via the EVLA Commissioning Staff Observing (ECSO) program.
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| == Expected Capabilities: Antennas ==
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| Retrofitted EVLA antennas are being returned to the array to be used as part of normal operations at the rate of approximately one antenna every two months. At the beginning of EVLA early science there will be 26 antennas in the array. The remaining two VLA antennas will be decommissioned while their retrofits are completed; they can not be used in conjunction with the EVLA correlator until they have been converted to the EVLA antenna design. All conversions will be completed by mid-2010.
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| == Expected Capabilities: Receivers ==
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| All retrofitted EVLA antennas are outfitted with either EVLA or “interim” L, EVLA or “interim” C, VLA X, EVLA K, and EVLA Q-band receivers. (Interim receivers are EVLA receivers with narrowband VLA polarizers. All interim receivers will be converted to full EVLA capabilities by the end of 2012. The polarization purity and sensitivity of the interim receivers typically is good only over the traditional VLA tuning range.) As of January 2010, 17 of the EVLA antennas are also outfitted with EVLA Ka-band receivers, and 6 EVLA antennas have S-band receivers. Figure 1 shows the expected rate of antenna retrofits and installation of the final EVLA receiver systems throughout the EVLA construction project. The 8-GHz maximum bandwidth availability depends on the implementation of the fast 3-bit samplers (the “8 GHz BW” line in Figure 1). Prior to this, the maximum available bandwidth will be 2 GHz per polarization.
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| [[image:EVLA_WidebandRxAvailability.png]] | |
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| '''Figure 1''': EVLA Receiver Deployment Plan. Above is a plot of the availability of the final EVLA receivers from late 2009 until the end of the EVLA Construction Project in 2012. Only final EVLA receivers are shown. Interim receivers with reduced frequency coverage or polarization purity are available at some bands (see Table 4). Approximate installation dates for the full 8 GHz bandwidth per polarization also are shown.
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| Figure 1 does not tell the entire story of frequency availability for observing with the EVLA, however, since there are interim or VLA receivers at L, C, and X-bands that can be used in the absence of the final EVLA receivers. Table 4 gives a prediction of the new frequency capabilities that we expect in June 2010, along with the expected “total” numbers of receivers for a given band, including VLA-style and/or interim receivers. New receiver bands will be offered for general use when the performance of at least five antennas has been verified by EVLA commissioning staff.
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| {| border="1"
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| |+ '''Table 4: Tuning Ranges of EVLA Bands'''
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| ! Band
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| ! Range
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| ! colspan="2" | Receiver availability, June 2010
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| |-
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| !
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| ! GHz
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| ! Final EVLA systems
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| ! Total EVLA+VLA/interim
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| |-
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| | 20 cm (L) || 1.0-2.0 || align='center'| 9 || align='center'| 26
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| |-
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| | 13 cm (S) || 2.0-4.0 || align='center'| 10 || align='center'| 10
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| |-
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| | 6 cm (C) || 4.0-8.0 || align='center'| 22 || align='center'| 26
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| |-
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| | 3 cm (X) || 8.0-12.0 || align='center'| 0 || align='center'| 27
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| | 2 cm (Ku) || 12.0-18.0 || align='center'|5 || align='center'| 5
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| |-
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| | 1.3 cm (K) || 18.0-26.5 || align='center'|27 || align='center'| 27
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| | 1 cm (Ka) || 26.5-40.0 || align='center'|25 || align='center'|25
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| |-
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| | 0.7 cm (Q) || 40.0-50.0 || align='center'|27 || align='center'|27
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| |-
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| |}
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| Note: The rightmost column gives the total numbers of receivers expected to be available for a given band, including all VLA-style and/or interim receivers. Only the old narrow-band VLA receivers are available in the 8-8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1).
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| = Performance of the EVLA =
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| This section contains details of the EVLA’s resolution, expected sensitivity, tuning range, dynamic range, pointing accuracy, and modes of operation. Detailed discussions of most of the observing limitations are found elsewhere. In particular, see References 1 and 2, listed in Section 7.
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| == Resolution ==
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| The EVLA’s resolution is generally diffraction-limited, and thus is set by the array config- uration 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 EVLA acts like any single antenna – structures smaller than the diffraction limit (θ ∼ λ/D) are broadened to the resolution of the antenna. 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.
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| 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), and the scale at which severe attenuation of large scale structure occurs.
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| A project with the goal of doubling the longest baseline available in the A configuration by establishing a real-time fiber optic link between the VLA and the VLBA antenna at Pie Town was established in the late 1990s, and used through 2005. This link is no longer operational; there is a goal (unfunded, at present) of implementing a new digital Pie Town link after the EVLA construction project has been completed.
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| == Sensitivity ==
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| The theoretical thermal noise expected for an image using natural weighting of the visibility
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| data is given by:
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| [[Image:thermalnoise.png]]
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| where:
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| :– ''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 EVLA’s 25-meter paraboloids, the ''SEFD'' is given by the equation ''SEFD'' = 5.62T<sub>sys</sub>/η<sub>A</sub>, where T<sub>sys</sub> is the total system temperature (receiver plus antenna plus sky), and η<sub>A</sub> is the antenna aperture efficiency in the given band.
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| :– η<sub>c</sub> is the correlator efficiency (at least 0.92 for the EVLA).
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| :– n<sub>pol</sub> is the number of polarization products included in the image; n<sub>pol</sub> = 2 for images in Stokes I, Q, U, or V, and npol = 1 for images in ‘RCP’ or ‘LCP’.
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| :– N is the number of antennas.
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| :– t<sub>int</sub> is the total on-source integration time in seconds.
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| :– ∆ν is the bandwidth in Hz.
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| Figure 2 shows the ''SEFD'' as a function of frequency for those bands currently installed on EVLA antennas, and include the contribution to Tsys from atmospheric emission at the zenith. Table 6 gives the SEFD at some fiducial EVLA frequencies. At X-band, where the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. Because EVLA testing with WIDAR has been used with a limited subset of antennas, it has not yet been possible to test whether equation 1 holds for an image made using the full array. However, it is expected that this will be the case, and equation 1 should be used for estimating required integration times.
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| {| border="1"
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| |+ '''Table 5: Configuration Properties'''
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| !Configuration
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| !A
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| !B
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| !C
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| !D
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| |-
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| |B<sub>max</sub> (km<sup>1</sup>) || 36.4 || 11.1 || 3.4 || 1.03
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| |-
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| |B<sub>min</sub> (km<sup>1</sup>) || 0.68 || 0.21 || 0.035<sup>5</sup> || 0.035
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| |-
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| | ||colspan="4"| Synthesized Beamwidth theta<sub>HPBW</sub>(arcsec)<sup>1,2,3</sup>
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| |-
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| |1.5 GHz (L) || 1.3 || 4.3 || 14 || 46
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| |-
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| |3.0 GHz (S)<sup>6</sup> || 0.65 || 2.1 || 7.0 || 23
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| |-
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| |6.0 GHz (C) || 0.33 || 1.0 || 3.5 || 12
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| |-
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| |8.5 GHz (X)<sup>7</sup> || 0.23 || 0.73 || 2.5 || 8.1
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| |-
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| |15 GHz (Ku)<sup>6</sup> || 0.13 || 0.42 || 1.4 || 4.6
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| |-
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| |22 GHz (K) || 0.089 || 0.28 || 0.95 || 3.1
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| |-
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| |33 GHz (Ka)<sup>6</sup> || 0.059 || 0.19 || 0.63 || 2.1
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| |-
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| |45 GHz (Q) || 0.043 || 0.14 || 0.47 || 1.5
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| |-
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| | || colspan="4" | Largest Angular Scale theta<sub>LAS</sub>(arcsec)<sup>1,4</sup>
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| |-
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| |1.5 GHz (L) || 36 || 120 || 970 || 970
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| |-
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| |3.0 GHz (S)<sup>6</sup> || 18 || 58 || 490 || 490
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| |-
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| |6.0 GHz (C) || 8.9 || 29 || 240 || 240
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| |-
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| |8.5 GHz (X)<sup>7</sup> || 6.3 || 20 || 170 || 170
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| |-
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| |15 GHz (Ku)<sup>6</sup> || 3.6 || 12 || 97 || 97
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| |-
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| |22 GHz (K) || 2.4 || 7.9 || 66 || 66
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| |-
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| |33 GHz (Ka)<sup>6</sup> || 1.6 || 5.3 || 44 || 44
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| |-
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| |45 GHz (Q) || 1.2 || 3.9 || 32 || 32
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| |-
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| |}
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| 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.
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| Footnotes:
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| :1. B<sub>max</sub> is the maximum antenna separation, B<sub>min</sub> is the minimum antenna separation, theta<sub>HPBW</sub> is the synthesized beam width (FWHM), and theta<sub>LAS</sub> is the largest scale structure "visible" to the array.
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| :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 (BnA, CnB, DnC) 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., CnA, or DnB) is provided.
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| :3. The approximate resolution for a naturally weighted map is about 1.5 times the numbers listed for theta<sub>HPBW</sub>. The values for snapshots are about 1.3 times the listed values.
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| :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.
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| :5. The standard '''C''' configuration has been replaced by a slightly modified one, formerly known as CS, wherein an antenna from the middle of the north arm has been moved to the central pad “N1”. This results in improved imaging for extended objects, but will degrade snapshot performance. Although the minimum spacing is the same as in '''D''' configuration, the surface brightness sensitivity to extended structure is considerably inferior to that of the '''D''' configuration (but considerably better than standard C configuration).
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| :6. The S, Ku, and Ka bands do not yet have a full complement of antennas, so the exact values will depend on the rate of antenna outfitting and the placement of individual antennas in the various configurations.
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| :7. At X-band the default VLA frequency of 8.5 GHz has been assumed, since there are no EVLA 8–12 GHz receivers available yet and the VLA-style receivers will probably be used through the end of 2010.
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| [[Image:SEFD.png|Figure 2: SEFD for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.]]
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| : '''Figure 2:''' ''SEFD'' for the EVLA. Left: The system equivalent flux density as a function of frequency for the L, S, and C-band receivers; note the logarithmic frequency axis. Right: The system equivalent flux density as a function of frequency for the K, Ka, and Q-band receivers. The frequency axis is linear.
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| {| border="1"
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| |+ '''Table 6: SEFDs and D-Configuration Confusion Limits'''
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| !Frequency
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| !SEFD
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| !RMS confusion level
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| |-
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| !
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| ! (Jy)
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| ! in '''D''' config (\mu Jy/beam)
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| |-
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| | 1.5 GHz (L) || 420 || 89
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| |-
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| | 3.0 GHz (S) || 370 || 14
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| |-
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| | 6.0 GHz (C) || 310 || 2.3
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| |-
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| | 22 GHz (K) || 560 || ...
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| |-
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| | 33 GHz (Ka) || 600 || ...
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| |-
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| | 45 GHz (Q) || 1400 || ...
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| |}
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| Note: SEFDs at K, Ka, and Q bands include contributions from Earth's atmosphere, and were determined under good conditions. At X-and, wehre the VLA-style receivers are still in use, the ''SEFD'' is approximately 310 Jy. The confusion limits in '''C''' configuration are approximately a factor of 10 less than those listed above.
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| 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:
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| :• 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.
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| :• 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.
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| :• Weather. The sky and ground temperature contributions to the total system temper- ature 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.
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| :In general, the zenith atmospheric opacity to microwave radiation is very low – typi- cally 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.
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| :Observers should remember that clouds, especially clouds with large water droplets (read, thunderstorms!), can add appreciable noise to the system temperature. Sig- nificant increases in system temperature can, in the worst conditions, be seen at frequencies as low as 5 GHz.
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| :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. 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.
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| | |
| :• Antenna elevation-dependent gains. The antenna figure degrades at low elevations, leading to diminished forward gain at the shorter wavelengths. The gain-elevation effect is negligible at frequencies below 8 GHz. The antenna gains can be determined by direct measurement of the relative system gain using the AIPS task ELINT on data from a strong calibrator which has been observed over a wide range of elevation. If this is not possible, care should be taken to observe a primary flux calibrator at the same elevation as the target.
| |
| | |
| :The AIPS task INDXR applies standard elevation-dependent gains and an estimated opacity in CL table version 1. The CASA calibration tasks (e.g. gaincal, bandpass) also use the standard gain curves. Note that currently the gain coefficients used in AIPS and CASA are for VLA antennas; the elevation-dependent gains of the EVLA antennas have yet to be fully characterized.
| |
| | |
| :• Pointing. The SEFD quoted assumes good pointing. Under calm nighttime condi- tions, the antenna blind pointing is about 10 arcsec rms. The pointing accuracy in daytime is a little worse, 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 if the reference source is within about 10 degrees (in azimuth and elevation) of the target source, and the source elevation is less than 70 degrees.
| |
| | |
| :Use of referenced pointing is highly recommended for all K, Ka, and Q-band observa- tions, 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 time. The calibrator should have at least 0.5 Jy flux density at X-band and be unresolved on all baselines to ensure an accurate solution.
| |
| | |
| To aid EVLA proposers there is an exposure tool calculator on-line at http://science.nrao.edu/evla/tools/exposure/evlaExpoCalc.jnlp that provides a graphical user interface to these equations.
| |
| | |
| 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:
| |
| | |
| [[Image:BrightnessTemp.png]] | |
| | |
| 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/B<sub>max</sub>), 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, ∆I<sub>m</sub>, 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 Section 7.
| |
| | |
| 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):
| |
| | |
| [[Image:MHI.png]]
| |
| | |
| where D is the distance to the galaxy in M<sub>pc</sub>, and S∆V is the Hi line area in units of Jy km/s.
| |
| | |
| == EVLA Frequency Bands and Tunability ==
| |
| | |
| For OSRO observations each receiver can tune to two different frequencies from the same wavelength band. Right-hand circular (RCP) and left-hand circular (LCP) polarizations are received for both frequencies. Each of these four data streams currently follows the VLA nomenclature, and are known as IF (for “Intermediate Frequency” channel) “A”, “B”, “C”, and “D”. IFs A and B receive RCP, IFs C and D receive LCP. IFs A and C are always at the same frequency, as are IFs B and D (but the IFs A/C frequency is usually different from the B/D frequency). We normally refer to these two independent data streams as “IF pairs.”
| |
| | |
| The tuning ranges, along with default frequencies for continuum applications, are given in Table 7 below. The EVLA X and Ku bands are not yet available, although the old narrow-band VLA X-band receivers may still be used.
| |
| | |
| {| border="1"
| |
| |+ '''Table 7: Default frequencies for "continuum" applications'''
| |
| ! Band
| |
| ! Range
| |
| ! colspan="2"| Default frequencies for continuum applications (GHz)
| |
| |-
| |
| !
| |
| ! (GHz)
| |
| ! IFs A/C
| |
| ! IFs B/D
| |
| |-
| |
| | 20 cm (L) || 1.0-2.0 || 1.328<sup>1</sup> || 1.456<sup>1</sup>
| |
| |-
| |
| | 13 cm (S) || 2.0-4.0 || 3.084 || 3.212
| |
| |-
| |
| | 6 cm (C) || 4.0-8.0 || 4.896 || 5.024
| |
| |-
| |
| | 3 cm (X) || 8.0-12.0 || 8.396<sup>2</sup> || 8.524<sup>2</sup>
| |
| |-
| |
| | 2 cm (Ku) || 12.0-18.0 || ... || ...
| |
| |-
| |
| | 1.3 cm (K) || 18.0-26.5 || 22.396 || 22.524
| |
| |-
| |
| | 1 cm (Ka) || 26.5-40.0 || 33.496 || 33.624
| |
| |-
| |
| | 0.7 cm (Q) || 40.0-50.0 || 43.216 || 43.344
| |
| |}
| |
| | |
| :'''Notes:'''
| |
| :: 1. This default frequency set-up for L-band avoids known interference signals as far as possible, but includes the Galactic HI line; another possibility is (a) to place both frequencies in the interference- free region at the high end of the band, at 1732 and 1860 MHz (see Figure 3), or (b) to separate the two frequencies placing one at 1452 MHz and the other at 1860 MHz, although this will make imaging the combined data more difficult due to the frequency difference between the IFs.
| |
| :: 2. Only the old narrow-band VLA receivers are available in the 8–8.8 GHz band until closer to the end of the EVLA construction project (see Figure 1). The default frequencies above are for the VLA X-band receivers; they will likely be changed for the final EVLA X-band receivers.
| |
| | |
| 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 Ka and Q bands (the only two 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.
| |
| :• At Q band, if the frequency span is greater than 8.0 GHz, the BD frequency must be lower than the AC frequency.
| |
| | |
| == Field of View ==
| |
| | |
| At least four different effects will limit the field of view. These are: primary beam; chromatic aberration; time-averaging; and non-coplanar baselines. We discuss each briefly:
| |
| | |
| === Primary Beam ===
| |
| | |
| The ultimate factor limiting the field of view is the diffraction-limited response of the indi- vidual antennas. An approximate formula for the full width at half power in arcminutes is: θ<sub>PB</sub> = 45/ν<sub>GHz</sub>. 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 for details.
| |
| | |
| === Chromatic Aberration (Bandwidth Smearing) ===
| |
| | |
| The principles upon which synthesis imaging are based are strictly valid only for monochro- matic 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 (∆ν/ν<sub>0</sub>) with the source offset in synthesized beamwidths (θ<sub>0</sub>/θ<sub>HPBW</sub>). 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.
| |
| | |
| {| border="1"
| |
| |+ '''Table 8: Reduction in Peak Response Due to Bandwidth Smearing'''
| |
| ! ∆ν/ν<sub>0</sub> θ<sub>0</sub>/θ<sub>HPBW</sub>
| |
| ! Peak
| |
| ! Width
| |
| |-
| |
| | 0.0 || 1.00 || 1.00
| |
| |-
| |
| | 0.50 || 0.95 || 1.05
| |
| |-
| |
| | 0.75 || 0.90 || 1.11
| |
| |-
| |
| | 1.0 || 0.80 || 1.25
| |
| |-
| |
| | 2.0 || 0.50 || 2.00
| |
| |}
| |
| | |
| :'''Note:''' The reduction in peak response and increase in width of an object due to bandwidth smearing (chromatic aberration). ∆ν/ν0 is the fractional bandwidth; θ0/θHPBW is the source offset from the phase tracking center in units of the synthesized beam.
| |
| | |
| === Time-Averaging Loss ===
| |
| | |
| The sampled coherence function (visibility) for objects not located at the phase-tracking center is slowly time-variable due to the motion of the source through the interferometer coherence pattern, so that averaging the samples in time will cause a loss of amplitude. Unlike the bandwidth loss effect described above, the losses due to time averaging cannot be simply parameterized, except for observations at δ = 90<sup>◦</sup>. 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 ∆ν/ν<sub>0</sub> replaced by ω<sub>e</sub>∆t<sub>int</sub>, where ω<sub>e</sub> is the Earth’s angular rotation rate, and ∆t<sub>int</sub> 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 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 (θ∆t<sub>int</sub>)<sup>2</sup>, where θ is the distance from the phase center and ∆t<sub>int</sub> 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 Section 4.5) in at least the '''A''' and '''B''' configurations.
| |
| | |
| {| border="1"
| |
| |+ '''Table 9: Averaging Time for a Given Amplitude Loss'''
| |
| !
| |
| ! colspan="3"|Amplitude loss
| |
| |-
| |
| ! Configuration
| |
| ! 1.0%
| |
| ! 5.0%
| |
| ! 10.0%
| |
| |-
| |
| | '''A''' || 2.1 || 4.8 || 6.7
| |
| |-
| |
| | '''B''' || 6.8 || 15.0 || 21.0
| |
| |-
| |
| | '''C''' || 21.0 || 48.0 || 67.0
| |
| |-
| |
| | '''D''' || 68.0 || 150.0 || 210.0
| |
| |}
| |
| | |
| :'''Note:''' The averaging time (in seconds) resulting in the listed amplitude losses for a point source at the antenna first null. Multiply the tabulated averaging times by 2.4 to get the amplitude loss at the half-power point of the primary beam. Divide the tabulated values by 4 if interested in the amplitude loss at the first null for the longest baselines.
| |
| | |
| === Non-Coplanar Baselines ===
| |
| | |
| The procedures by which nearly all images are made in Fourier synthesis imaging are based on the assumption that all the coherence measurements are made in a plane. This is strictly true for E-W interferometers, but is false for the EVLA, with the single exception of snapshots. Analysis of the problem shows that the errors associated with the assumption of a planar array increase quadratically with angle from the phase-tracking center. Serious errors result if the product of the angular offset in radians times the angular offset in synthesized beams exceeds unity: θ > λB/D2, where B is the baseline length, D is the antenna diameter, and λ is the wavelength, all in the same units. This effect is most noticeable at λ90 and λ20 cm in the larger configurations, but will be notable in wide-field, high fidelity imaging for other bands and configurations.
| |
| | |
| Solutions to the problem of imaging wide-field data taken with non-coplanar arrays are well known, and have been implemented in AIPS (IMAGR) and CASA (clean). Refer to the package help files for these tasks, or consult with Rick Perley, Frazer Owen, or Sanjay Bhatnagar for advice. More computationally efficient imaging with non-coplanar baselines is being investigated, such as the “W-projection” method available in CASA; see EVLA Memo 67 for more details.
| |
| | |
| == Time Resolution and Data Rates ==
| |
| | |
| The minimum integration time that will be supported for OSRO is 1 second. It is expected, however, that for normal observing in the D configuration integration times of up to 10 seconds will be commonly used.
| |
| | |
| The maximum recommended integration time for any EVLA observing is 60 seconds, but given current computing capabilities there is no reason to use integration times longer
| |
| than 10 seconds. For high frequency observers with short scans (e.g., fast switching, as described in Section 4.11.2), shorter integration times may be preferable.
| |
| | |
| Observers should bear in mind the data rate of the EVLA when planning their observa- tions. For the correlator configurations discussed in Section 4.13, and integration time ∆t, the data rate is:
| |
| | |
| :Data rate = 1.7 MB/sec×Nant ×(Nant −1)/(27×26)/(∆t/1 sec)
| |
| :::= 6.0 GB/hour×Nant ×(Nant −1)/(27×26)/(∆t/1 sec) (4)
| |
| | |
| which corresponds to 145 GB/day when using the full 27-element array with 1 second integrations. While analyzing such data sets is not too difficult with current computers, data transfer will clearly be more of an issue than in the past. Most observers currently download their data via ftp directly from the archive. The Archive Access Tool will allow 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. Observers may also request their data on DATs or USB drives.
| |
| | |
| == Radio-Frequency Interference ==
| |
| | |
| The bands within the tuning range of the EVLA 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 is primarily a problem within the low frequency bands, and is most serious to the D configuration, as the fringe rates in other configurations are often sufficient to reduce interference to tolerable levels.
| |
| | |
| Radio frequency interference (RFI) at the EVLA will be an increasing problem to as- tronomical observations. Table 10 lists some of the sources of external RFI at the VLA site that might be observed within the EVLA tuning range. Figure 3 shows a raw power spec- trum, using the prototype WIDAR0 correlator. Figure 4 shows a similar plot for the new S-band system. Table 10 gives a summary of the origins (where known) of the prominent features shown in the two figures.
| |
| | |
| Plots of all RFI observations from 1993 onwards are available online, at http://www.vla.nrao.edu/cgi-bin/rfi.cgi. For general information about the RFI en- vironment, contact the head of the IPG (Interference Protection Group) by sending e-mail to nrao-rfi@nrao.edu.
| |
| | |
| The impact of RFI on astronomical observing depends on the configuration and wave- length. It is worst for the D configuration, but RFI effects, as seen in an image, will be appreciably attenuated in the extended configurations and at higher frequencies due to fringe phase winding.
| |
| | |
| The EVLA electronics (including the WIDAR correlator) have been designed to mini- mize gain compression due to very strong RFI signals, so that in general it will be possible to observe in spectral regions containing RFI, provided the spectra are well sampled to minimize Gibbs ringing, and spectral smoothing (such as Hanning) is applied. We fully expect useful astronomical data to be extracted even if extremely strong interfering signals are located a few MHz away.
| |
| | |
| Extracting astonomy 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.
| |
| | |
| {| border="1"
| |
| |+ '''Table 10: Identified VLA RFI Between 1 and 4 GHz'''
| |
| ! Frequency (MHz)
| |
| ! Source
| |
| ! Comments
| |
| |-
| |
| | 1025-1150 || Aircraft navigation || Very strong
| |
| |-
| |
| | 1200.0 || VLA modem ||
| |
| |-
| |
| | 1217-1237 || GPS L2 || Very strong
| |
| |-
| |
| | 1243-1251 || GLONASS L2 ||
| |
| |-
| |
| | 1254 || Aeronautical radar ||
| |
| |-
| |
| | 1263 || Aeronautical radar ||
| |
| |-
| |
| | 1268 || COMPASS E6 ||
| |
| |-
| |
| | 1310 || Aeronautical radar ||
| |
| |-
| |
| | 1317 || Aeronautical radar ||
| |
| |-
| |
| | 1330 || Aeronautical radar ||
| |
| |-
| |
| | 1337 || Aeronautical radar ||
| |
| |-
| |
| | 1376-1386 || GPS L3 || Intermittent
| |
| |-
| |
| | 1525-1564 || INMARSAT satellites ||
| |
| |-
| |
| | 1564-1584 || GPS L1 || Very strong
| |
| |-
| |
| | 1598-1609 || GLONASS L1 ||
| |
| |-
| |
| | 1618-1627 || IRIDIUM satellites ||
| |
| |-
| |
| | 1642 || 2nd harmonic VLA radios || Sporadic
| |
| |-
| |
| | 1683-1687 || GOES weather satellite ||
| |
| |-
| |
| | 1689-1693 || GOES weather satellite ||
| |
| |-
| |
| | 1700-1702 || NOAA weather satellite ||
| |
| |-
| |
| | 1705-1709 || NOAA weather satellite ||
| |
| |-
| |
| | 1930-1990 || PCS cell phone base stations ||
| |
| |-
| |
| | 2178-2195 || ??? ||
| |
| |-
| |
| | 2320-2350 || Satellite radio ||
| |
| |-
| |
| |}
| |
| | |
| == Subarrays ==
| |
| | |
| The separation of the EVLA into multiple sub-arrays will not be supported initially for OSRO.
| |
| | |
| == Positional Accuracy ==
| |
| | |
| The accuracy with which an object’s position can be determined is limited by the at- mospheric phase stability, the closeness of a suitable (astrometric) calibrator, and the calibrator-source cycle time. Under good conditions, in A configuration, accuracies of about 0.05 arcseconds can be obtained. Under more normal conditions, accuracies of perhaps 0.1 arcseconds can be expected. Under extraordinary conditions (probably attained only a few times per year on calm winter nights in A configuration when using rapid phase switching on a nearby astrometric calibrator – see Section 4.11.2), accuracies of 1 milliarcsecond have been attained with the VLA.
| |
| | |
| If highly accurate positions are desired, only “A” code (astrometric) calibrators from the VLA Calibrator List should be used. The positions of these sources are taken from lists published by the USNO.
| |
| | |
| [[Image:Lband_sweep.png|Figure 3: Spectrum of L-band RFI. This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.]]
| |
| | |
| :'''Figure 3: Spectrum of L-band RFI.''' This shows the major interfering signals seen across the full 1 GHz bandwidth available to the L-band receivers. Each of the eight “spectral windows” displays 128 MHz from a separate sub-band. These are raw data, uncalibrated for the bandpass of either the digital filter or the receiver. The high linearity of the EVLA’s electronics and correlator will permit astronomical observing within any frequencies not containing external interference.
| |
| | |
| == Limitations on Imaging Performance ==
| |
| | |
| Imaging performance can be limited in many different ways. Some of the most common are listed in the following subsections.
| |
| | |
| === Image Fidelity ===
| |
| | |
| With conventional point-source calibration methods, and even under the best observing conditions, the achieved dynamic range will rarely exceed a few hundred. The limiting factor is usually the atmospheric phase stability. If the target source contains more than 50 mJy in compact structures (depending somewhat on band), self-calibration can be counted on to improve the images. If the atmospheric coherence time is several minutes or longer, weaker sources can be used for self-cal; for shorter coherence times and less sensitive bands, stronger sources will be needed. Dynamic ranges in the thousands can be achieved using these techniques. With the new WIDAR correlator, much greater bandwidths and much higher sensitivities are available, and we expect self-calibration methods will be extendable to observations of sources with much lower flux densities than the current limits.
| |
| | |
| [[Image:S-bandRFI.png|This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.]]
| |
| | |
| :'''Figure 4: Spectrum of S-band RFI.''' This shows the raw spectrum of the lower half of S- Band – 2.0 to 3.0 GHz. No significant RFI is seen in the upper half. The major interference at 2.35 GHz is from satellite radio.
| |
| | |
| === Invisible Structures ===
| |
| | |
| An interferometric array acts as a spatial filter, so that for any given configuration, struc- tures 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. Table 5 gives the largest scale visible to each configuration/band combination.
| |
| | |
| === Poorly Sampled Fourier Plane ===
| |
| | |
| Unmeasured Fourier components are assigned values by the deconvolution algorithm. While this often works well, sometimes it fails noticeably. The symptoms depend upon the actual deconvolution algorithm used. For the CLEAN algorithm, the tell-tale sign is a fine mottling on the scale of the synthesized beam, which sometimes even organizes itself into coherent stripes. Further details are to be found in Reference 1.
| |
| | |
| === Sidelobes from Confusing Sources ===
| |
| | |
| At the lower frequencies, large numbers of detectable background sources are located throughout the primary antenna beam, and into its first sidelobe. Sidelobes from those sources which have not been deconvolved will lower the image quality of the target source. Although bandwidth 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.
| |
| | |
| === 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.
| |
| | |
| == Calibration and the Flux Density Scale ==
| |
| | |
| The VLA Calibrator List contains information on 1860 sources sufficiently un- resolved and bright to permit their use as calibrators. The list is available within the Observation Preparation Tool and may be accessed on the Web at http://www.vla.nrao.edu/astro/calib/manual/.
| |
| | |
| Accurate flux densities can be obtained by observing one of 3C286, 3C147, 3C48 or 3C138 during the observing run. Not all of these are suitable for every observing band and configuration – consult the VLA Calibrator Manual for advice. Over the last several years, we have implemented accurate source models directly in AIPS and CASA for much improved calibration of the amplitude scales. Models are available for 3C48, 3C138, 3C147, and 3C286 for L, C, X, Ku, K, and Q bands. At Ka band either of the K or Q band models works reasonably well.
| |
| | |
| Since the standard source flux densities are slowly variable, we monitor their flux den- sities when the array is in its D configuration. As the EVLA cannot 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 November 2001 at the standard VLA bands. The accuracy of these values, relative to the assumed standards, is set by the gain stability of the instrument. The estimated 1-σ errors in the table, relative to the assumed standards, are less than 1% for frequencies up to 25 GHz, and about 2% for the 43 GHz band.
| |
| | |
| The flux densities shown in the table for frequencies below 10 GHz are based on the Baars et al. value for 3C286. For frequencies above 10 GHz, the flux densities are based on a model of the emission of Mars.
| |
| | |
| Polynomial coefficients describing the derived flux densities for the standard calibrators have been determined which permit accurate interpolation of the flux density at any EVLA frequency. These coefficients are updated approximately every few years, and are used in the AIPS task SETJY and in the CASA task setjy. A substantial effort is under way to establish the long-term (past and present) accuracy of the EVLA flux density scale; contact Rick Perley or Bryan Butler for information on this work. Although final information is not yet available, it is clear that the flux density of 3C286 has not changed by more than 1% over the last 25–30 years at any band – i.e., the flux density of 3C286 appears to be stable.
| |
| | |
| {| border="1"
| |
| |+ '''Table 11: Flux densities (Jy) of Standard Calibrators for November 2001'''
| |
| ! Source/Frequency (MHz)
| |
| ! 1465
| |
| ! 4885
| |
| ! 8435
| |
| ! 14965
| |
| ! 22460
| |
| ! 43340
| |
| |-
| |
| | 3C48 = J0137+3309 || 15.54 || 5.47 || 3.22 || 1.83 || 1.23 || 0.63
| |
| |-
| |
| | 3C138 = J0521+1638 || 8.16 || 3.71 || 2.43 || 1.54 || 1.10 || 0.62
| |
| |-
| |
| | 3C147 = J0542+4951 || 21.32 || 7.88 || 4.72 || 2.77 || 1.93 || 1.11
| |
| |-
| |
| | 3C286 = J1331+3030 || 14.49 || 7.49 || 5.22 || 3.51 || 2.63 || 1.58
| |
| |-
| |
| | 3C295 = J1411+5212 || 21.49 || 6.53 || 3.42 || 1.67 || 0.99 || 0.41
| |
| |-
| |
| | NGC 7027 || 1.54 || 5.52 || 6.03 || 5.90 || 5.68 || 5.22
| |
| |-
| |
| | MARS || - || 0.175 || 0.528 || 1.67 || 3.81 || 14.22
| |
| |-
| |
| |}
| |
| | |
| For most observing projects, the effects of atmospheric extinction will automatically be accounted for by regular calibration when using a nearby point source whose flux density has been determined by an observation of a flux density standard taken at a similar elevation. However, at high frequencies (i.e., K-band, Ka-band, and Q-band), both the variation of antenna gain and the atmospheric absorption with elevation may be strong enough to make “simple” flux density bootstrapping unreliable. The AIPS task ELINT is available to permit measurement of an elevation gain curve using your own observations, and subsequent adjustment of the derived gains to remove these elevation-dependent effects. The current calibration methodology does not require knowledge of the atmospheric extinction (since the true flux densities of the standard calibrators are believed known). However, if knowledge of the actual extinction is desired, tipping scans can be included in an observation.
| |
| | |
| == Complex Gain Calibration ==
| |
| | |
| === General Guidelines for Gain Calibration ===
| |
| | |
| Adequate gain calibration is a complicated function of source-calibrator separation, fre- quency, array scale, and weather. And, since what defines adequate for some experiments is completely inadequate for others, it is impossible to define any simple guidelines to ensure adequate phase calibration in general. However, some general statements remain valid most of the time. These are given below.
| |
| | |
| :• Tropospheric effects dominate at wavelengths shorter than 20 cm, ionospheric effects dominate at wavelengths longer than 20 cm.
| |
| :• Atmospheric (troposphere and ionosphere) effects are nearly always unimportant in the '''C''' and '''D''' configurations at L and S bands, and in the '''D''' configuration at X and C bands. Hence, for these cases, calibration need only be done to track instrumental changes – once per hour is generally sufficient.
| |
| :• If your target object has sufficient flux density to permit phase self-calibration, there is no need to calibrate more than once hourly at low frequencies (L/S/C bands) or 15 minutes at high frequencies (K/Ka/Q bands) in order to track pointing or other effects that might influence the amplitude scale.
| |
| :• The smaller the source-calibrator angular separation, the better. In deciding between a nearby “S” calibrator, and a more distant “P” calibrator, the nearby calibrator is usually the better choice.
| |
| :• At high frequencies, and longer configurations, rapid switching between the source and nearby calibrator is often helpful. See the following section.
| |
| | |
| === Rapid Phase Calibration and the Atmospheric Phase Interferometer (API) ===
| |
| | |
| For some objects, and under suitable weather conditions, the phase calibration can be considerably improved by rapidly switching between the source and calibrator. Source- Calibrator observing cycles as short as 40 seconds can be used. However, observing efficiency declines for very short cycle times, so it is important to balance this loss against a realistic estimate of the possible gain. Experience has shown that cycle times of 100 to 150 seconds at high frequencies have been effective for source-calibrator separations of less than 10 degrees. For the VLA this was known as “fast-switching.” For the EVLA it is just a loop of source-calibrator scans with short scan length. This technique “stops” tropospheric phase variations at an effective baseline length of ∼ vat/2 where va is the atmospheric wind velocity aloft (typically 10 to 15 m/sec), and t is the total switching time. It has been demonstrated to result in images of faint sources with diffraction-limited spatial resolution on the longest EVLA baselines. Under average weather conditions, and using a 120 second cycle time, the residual phase at 43 GHz should be reduced to ≤ 30 degrees. Further details can be found in VLA Scientific Memos # 169 and 173. These memos, and other useful information, can be obtained from Reference 12 in Section 7.
| |
| | |
| Note that the fast switching technique will not work in bad weather (such as rain showers, or when there are well-developed convection cells – most notably, thunderstorms).
| |
| | |
| An Atmospheric Phase Interferometer (API) is used to continuously measure the tro- pospheric 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 geosta- tionary 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 ob- server that high frequency observing may not be possible with current weather conditions. A detailed description of the API, and instructions for accessing its data, may be found at http://www.vla.nrao.edu/astro/guides/api/.
| |
| | |
| == 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 paral- lactic angle. In nearly all cases, the phase calibrator chosen can double as a polarization cali- brator. The minimum condition that will enable accurate polarization calibration is four ob- servations of a bright source spanning at least 90 degrees in parallactic angle. 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 to fix the position angle at wave- lengths of 6 cm or shorter. The results of a careful monitoring program of these and other polarization calibrators will be found at http://www.vla.nrao.edu/astro/calib/polar/.
| |
| | |
| High sensitivity linear polarization imaging may be limited by time dependent instru- mental 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 polarizes indicates that these are extremely stable over the duration of any single observation, strongly suggesting that high quality polarimetry over the full bandwidth will be possible.
| |
| | |
| The accuracy of wide field linear polarization imaging will be limited, likely at the level of a few percent at the antenna half-power width, by angular variations in the antenna polarization response. Algorithms to enable removable of this angle-dependent polarization are being tested, and observations to determine the antenna polarizations have begun.
| |
| | |
| Circular polarization measurements will be limited by the beam squint, due to the offset secondary focus feeds, which separates the RCP and LCP beams by a few percent of the FWHM. The same algorithms noted above to correct for antenna-induced linear polarization can be applied to correct for the circular beam squint. Measurement of the beam squints, and testing of the algorithms, is ongoing.
| |
| | |
| Ionospheric Faraday rotation of the astronomical signal is always notable at 20 cm. The typical daily maximum rotation measure under quiet solar conditions is 1 or 2 radians/m2, so the ionospherically-induced rotation of the plane of polarization at these bands is not excessive – 5 degrees at 20 cm. However, under active conditions, this rotation can be many times larger, sufficiently large that polarimetry is impossible at 20 cm with corrrection for this effect. The AIPS program TECOR has been shown to be quite effective in removing large-scale ionospherically induced Faraday Rotation. It uses currently-available data in IONEX format. Please consult the TECOR help file for detailed information.
| |
| | |
| In addition, the interim EVLA receivers generally have poor polarization performance outside the frequency range previously covered by the VLA (e.g., outside the 4.5–5.0 GHz frequency range for C band, and outside 1.3 – 1.7 GHz for L-band), and the wider frequency bands of these interim receivers may be useful only for total intensity measurements.
| |
| | |
| == Correlator Configurations ==
| |
| | |
| The EVLA correlator is very flexible, and will be able to provide data in many ways. Initially the EVLA correlator will be configurable for Open Shared Risk Observing in the following ways:
| |
| | |
| :1. “OSRO1”: A configuration for continuum applications, to comprise two independently tunable sub-bands, full polarization, each of which has 128 MHz bandwidth with 64 channels, with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. In this configuration it will also be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the spectro-polarimetry capabilities in Table 12. Both sub-bands must have the same bandwidth. This configuration utilizes both of the IF pairs (see Section 4.3).
| |
| | |
| :2. “OSRO2”: A higher spectral resolution configuration for spectral line applications, to comprise one tunable sub-band, providing dual polarization, with 128 MHz bandwidth and 256 channels, again with the possibility of smoothing in frequency to reduce dataset sizes or to improve spectral response offline. It will be possible to decrease the bandwidth by powers of two, keeping the same number of channels, to provide the capabilities in Table 13. This configuration utilizes either IF pair (see Section 4.3).
| |
| | |
| {| border="1"
| |
| |+ '''Table 12: Correlator capabilities per sub-band for full polarization continuum applications and spectro-polarimetry (OSRO1)'''
| |
| ! Sub-band BW (MHz)
| |
| ! Number of poln. products
| |
| ! Number of channels/poln. product
| |
| ! Channel width (kHz)
| |
| ! Channel width (km/s at 1 GHz)
| |
| ! Total velocity coverage per sub-band (km/s at 1 GHz)
| |
| |-
| |
| | 128 || 4 || 64 || 2000 || 600/nu (GHz) || 38,400/nu (GHz)
| |
| |-
| |
| | 64 || 4 || 64 || 1000 || 300 || 19,200
| |
| |-
| |
| | 32 || 4 || 64 || 500 || 150 || 9,600
| |
| |-
| |
| | 16 || 4 || 64 || 250 || 75 || 4,800
| |
| |-
| |
| | 8 || 4 || 64 || 125 || 37.5 || 2,400
| |
| |-
| |
| | 4 || 4 || 64 || 62.5 || 19 || 1,200
| |
| |-
| |
| | 2 || 4 || 64 || 31.25 || 9.4 || 600
| |
| |-
| |
| | 1 || 4 || 64 || 15.625 || 4.7 || 300
| |
| |-
| |
| | 0.5 || 4 || 64 || 7.813 || 2.3 || 150
| |
| |-
| |
| | 0.25 || 4 || 64 || 3.906 || 1.2 || 75
| |
| |-
| |
| | 0.125 || 4 || 64 || 1.953 || 0.59 || 37.5
| |
| |-
| |
| | 0.0625 || 4 || 64 || 0.977 || 0.29 || 18.75
| |
| |-
| |
| | 0.03125 || 4 || 64 || 0.488 || 0.15 || 9.375
| |
| |}
| |
| | |
| {| border="1"
| |
| |+ '''Table 13: Correlator capabilities per sub-band for spectral line applications, dual polarization (OSRO2)'''
| |
| ! Sub-band BW (MHz)
| |
| ! Number of poln. products
| |
| ! Number of channels/poln. product
| |
| ! Channel width (kHz)
| |
| ! Channel width (km/s at 1 GHz)
| |
| ! Total velocity coverage per sub-band (km/s at 1 GHz)
| |
| |-
| |
| | 128 || 2 || 256 || 500 || 150/nu (GHz) || 38,400/nu (GHz)
| |
| |-
| |
| | 64 || 2 || 256 || 250 || 75 || 19,200
| |
| |-
| |
| | 32 || 2 || 256 || 125 || 37.5 || 9,600
| |
| |-
| |
| | 16 || 2 || 256 || 62.5 || 19 || 4,800
| |
| |-
| |
| | 8 || 2 || 256 || 31.25 || 9.4 || 2,400
| |
| |-
| |
| | 4 || 2 || 256 || 15.625 || 4.7 || 1,200
| |
| |-
| |
| | 2 || 2 || 256 || 7.813 || 2.3 || 600
| |
| |-
| |
| | 1 || 2 || 256 || 3.906 || 1.2 || 300
| |
| |-
| |
| | 0.5 || 2 || 256 || 1.953 || 0.59 || 150
| |
| |-
| |
| | 0.25 || 2 || 256 || 0.977 || 0.29 || 75
| |
| |-
| |
| | 0.125 || 2 || 256 || 0.488 || 0.15 || 37.5
| |
| |-
| |
| | 0.0625 || 2 || 256 || 0.244 || 0.073 || 18.75
| |
| |-
| |
| | 0.03125 || 2 || 256 || 0.122 || 0.037 || 9.375
| |
| |}
| |
| | |
| These capabilities will be provided with integration times no shorter than 1 second, and it is hoped that Doppler tracking will be available with these correlator configurations.
| |
| | |
| The capabilities available to the OSRO program will be expanded as soon as they can be supported for general use. It is currently hoped that 2 GHz of bandwidth will be available for OSRO by the time of the D configuration beginning in Trimester 2, 2011.
| |
| | |
| If Doppler tracking is not used for an observation, or if it is likely that the data will need to be resampled spectrally in order to Doppler track to a different line rest frequency, care should be taken to make sure the spectrum is oversampled to avoid the subsequent introduction of Gibbs ringing or the need to reduce the spectral resolution by Hanning smoothing.
| |
| | |
| All observations with the EVLA correlator should be treated as traditional VLA spectral line observations, in that they will require observation of a bandpass calibrator. They may also require observation of a delay calibrator. Users should contact NRAO staff for advice on setting up observations with the EVLA correlator.
| |
| | |
| == VLBI Observations ==
| |
| | |
| VLBI observations with the EVLA, such as phased array and single-antenna modes, will not initially be available for the OSRO program.
| |
| | |
| == Snapshots ==
| |
| | |
| The two-dimensional geometry of the EVLA allows a snapshot mode whereby short obser- vations 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 FIRST or CORNISH ('''B''' configuration), or NVSS ('''D''' config- uration, all-sky) surveys. These surveys can be accessed from the NRAO website, at: http://science.nrao.edu/evla/proposing/largeproposals.shtml.
| |
| | |
| == Shadowing and Cross-Talk ==
| |
| | |
| Observations at low elevation in the '''C''' and '''D''' configurations will commonly be affected by shadowing. It is strongly recommended that all data from a shadowed antenna be discarded. This will automatically be done during filling (CASA task importasdm) when using the default inputs (Note: For archival VLA data, the AIPS task FILLM can also flag based on shadowing). AIPS task UVFLG can be used to flag 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. At 20 cm, this effect is important principally in the '''D''' configuration. Careful editing is necessary to identify and remove this form of interference.
| |
| | |
| == Combining Configurations and Mosaicing ==
| |
| | |
| Any single EVLA configuration will allow accurate imaging up to a scale approximately 30 times the synthesized beam. Objects larger than this will require multiple configuration observations. It is advisable that the frequencies used be the same for all configurations to be combined. Objects larger than the primary antenna pattern may be mapped through the technique of interferometric mosaicing. Time-variable structures (such as the nuclei of radio galaxies and quasars) cause special, but manageable, problems. See the article by Mark Holdaway in reference 2 for more information.
| |
| | |
| == Pulsar Observing ==
| |
| | |
| Observation of pulsars requiring modes other than those described in Section 4.13 will not initially be supported on the EVLA.
| |
| | |
| = Using the EVLA =
| |
| | |
| == Obtaining Observing Time on the EVLA ==
| |
| | |
| Observing time on the EVLA is available to all researchers, regardless of nationality or location of institution. There are no quotas or reserved blocks of time. The allocation of observing time on the EVLA is based upon the submission of an EVLA Observing Proposal using the on-line Proposal Submission Tool available via the NRAO Interactive Service web page, at http://my.nrao.edu/. The on-line tool permits the detailed construction of a cover sheet specifying the requested observations, using a set of on-line forms, and uploading of a pdf-format scientific and technical justification to accompany the cover information.
| |
| | |
| It is also possible to receive EVLA observing time by proposing to NASA missions, under cooperation agreements established between NRAO and those missions. Currently, such programs exist for the Chandra, Spitzer, and Fermi missions. Astronomers interested in those joint programs should consult the relevant mission proposal calls for more information.
| |
| | |
| Students planning to use the EVLA for their Ph.D. dissertation may have a problem in that such dissertations are frequently composed of pieces of several short proposals which may not be suitable for combining into a single proposal for refereeing purposes. In this case, we shall accept, one per student, a “Plan of Dissertation Research,” of no more than 1000 words, at the time of the first proposal of the series, and which can be referred to in later proposals. The plan can be submitted via the NRAO Interactive Services webpage, at http://my.nrao.edu/. This provides some assurance against a dissertation being seriously damaged by adverse referee comments on one component proposal, when the referees may not see the whole picture. This facility is offered to students for which EVLA observations are the most important component of their planned dissertations.
| |
| | |
| The EVLA is currently scheduled on a trimester basis, with each trimester lasting four months, matching the configuration cycle. The proposal deadline for a particular configura- tion is 5PM (1700), Eastern Time on the 1st of February, June, or October which precedes the beginning of that configuration by three months or more. (If the first of the month falls on a Saturday or Sunday, the deadline is advanced to the next Monday.) It is not necessary to submit a proposal at the last possible deadline for a particular configuration, since all proposals will be refereed immediately following the deadline of submission, regardless of the configuration requested. Early submissions (i.e., more than one deadline in advance of the relevant configuration deadline) will benefit from referee feedback and the opportunity for revision and additional review, if warranted.
| |
| | |
| All proposals are externally refereed by several experts in relevant subdisciplines (e.g., solar system, stellar, galactic, extragalactic, etc.). The referees’ comments and rating are strongly advisory to the VLA/VLBA Proposal Selection Committee, and the com- ments of both groups are passed on to the proposers soon after each meeting of the committee (three times yearly) and prior to the next proposal submission deadline. See http://science.nrao.edu/evla/proposing/timeallocation.shtml for a detailed description of the time allocation process.
| |
| | |
| Because of competition, even highly rated proposals are not guaranteed to receive ob- serving time. This is particularly true for programs that concentrate on objects in the LST ranges occupied by popular targets such as the Galactic Center or the Virgo cluster. Daytime observing will also be limited by EVLA commissioning throughout 2010.
| |
| | |
| == Rapid Response Science ==
| |
| | |
| The NRAO has established three categories of proposals for Rapid Response Science, which are described below and at http://science.nrao.edu/evla/proposing/rapid.shtml. At present, Rapid Response Science is limited to a maximum of 5% of the total observing time on the EVLA. All proposals for Rapid Response Science must be submitted using the standard NRAO procedures, using the on-line proposal tool. Proposals submitted by any other means (e.g., phone calls, e-mails, faxes, word-of-mouth) will be not be considered.
| |
| | |
| :1. '''Known Transient Phenomena.''' These proposals will request time to observe phe- nomena that are predictable in general, but not in specific detail. For example, a proposal to observe the next flaring X-ray binary that meets certain criteria would be included in this category. Specific triggering criteria will be required. These proposals will be evaluated as part of the normal refereeing and scheduling process, and will be subject to the normal NRAO proposal deadlines. The proprietary period for observations of Known Transient Phenomena will be 12 months.
| |
| :2. '''Exploratory Time.''' These proposals generally follow up on recent discoveries, with observations requested in advance of the period allocated at the most recent proposal dead- line(s). Examples include A configuration proposals that follow up on B configuration discoveries made after the A configuration deadline, or a newly identified object that is a hot enough topic to warrant an image within a couple of months. In general, there will not be a need for immediate scheduling with these proposals, but they may need to be observed in the current VLA configuration rather than waiting 16 months. Proposals for Exploratory Time will be evaluated by a subset of the EVLA/VLBA Proposal Selection Committee, and may or may not be sent to external referees. The possibility that a proposer forgets about or misses a proposal deadline will not constitute sufficient justification for granting of ob- serving time by this process. Notification of the dispensation of Exploratory Time proposals normally will be within two weeks of reception of the proposal; most of the accepted pro- posals are placed in the dynamic scheduling queue, and are not guaranteed to be observed. The proprietary period for Exploratory Time will be six months.
| |
| :3. '''Target of Opportunity.''' These proposals are for true targets of opportunity– unexpected or unpredicted phenomena such as supernovae, novae, or extreme X-ray, optical, or radio flares in various types of objects. These proposals will be evaluated rapidly, with scheduling done as quickly as possible and as warranted by the nature of the transient phenomenon. Notification of the dispensation of Target of Opportunity proposals will always be within two weeks, and may be much faster, depending on the requirements of the pro- posed observation. The proprietary period for Targets of Opportunity will be decided on a case-by-case basis, and will in no instances be longer than six months.
| |
| | |
| == Data Analysts and General Assistance ==
| |
| | |
| General assistance of all kinds is available through the data analysts, and they should be consulted first when you encounter any problem. Note that they are not available to perform remote data calibration. There are plans to provide pipeline-calibrated visibility data for the EVLA, but this will not be available initially. The e-mail address for all the data analysts is analysts@nrao.edu. An online helpdesk is in the process of being set up for assisting with EVLA observation preparation and data reduction. Specific questions about observation preparation and data reduction may also be directed to NRAO scientific staff via vlahelp@nrao.edu.
| |
| | |
| == Observation Preparation ==
| |
| | |
| To use the EVLA, scheduling blocks must be prepared using the
| |
| “Observation Preparation Tool,” or OPT. The OPT is available at: http://science.nrao.edu/evla/observing/opt.shtml.
| |
| | |
| == Fixed Date and Dynamic Scheduling ==
| |
| | |
| Most of the projects on the EVLA will be observed dynamically, based on a combination of scientific priority and the expected properties of the array and the weather. Some time may continue to be scheduled as “fixed-date” observing, with the observer being given a particular sidereal date and time allocation on the EVLA, depending on the needs of the project.
| |
| | |
| Please see the scheduling officers’ home page for further information, at http://www.aoc.nrao.edu/∼schedsoc/.
| |
| | |
| == The Observations and Remote Observing ==
| |
| | |
| Since most EVLA observations will be carried out dynamically it is in general neither necessary, nor practical, for observers to be present during their observations. However, there can be considerable benefits to observers who come to Socorro to reduce and analyze their data in terms of interactions and discussions with NRAO staff. See Section 5.12 for information on coming to and staying in Socorro.
| |
| | |
| For those who choose to process their data at home, the data can be retrieved from the EVLA online archive after obtaining the project key from the data analysts, or by using the NRAO Interactive Services login from the archive web page. Alternatively, the data analysts will, upon request, mail you a tape or other media containing your uncalibrated data in its original format.
| |
| | |
| == Data Access ==
| |
| The online archive contains all VLA data since observing started in 1976, and will also serve the user community with EVLA data. The entire archive is now on disk, and is available via the Archive Access Tool at http://archive.nrao.edu/archive/. This interface provides a basic data retrieval tool if you know the program code of your observations. It also provides an advanced query tool which enables searches based on a large number of user- specified criteria. Data can be downloaded via standard ftp protocols. With the exception of some rapid response and large proposals, (E)VLA data associated with a given proposal normally are restricted to proprietary use by the proposing team for a period of 12 months from the date of the last observation in a proposal (Note: Data taken more than 12 months previously may still be proprietary, if additional data for the same proposal have been taken within the last 12 months). Proprietary data may be downloaded by the observing team by making use of the project key (see Section 5.6).
| |
| | |
| Data are stored in the archive in the Science Data Model (SDM) format that will be used by both the EVLA and ALMA. They are available through the Archive Access Tool in one of several formats:
| |
| | |
| :– SDM format, which must be read into CASA using the task importasdm for further processing.
| |
| :– As a CASA Measurement Set (ms), which has already been converted from the SDM using default parameters in the CASA task importasdm.
| |
| :– UVFITS format, which as been converted from the SDM using CASA and then written to a multi-source UVFITS file using the CASA task exportuvfits, using default parameters. These UVFITS files can subsequently be read by AIPS.
| |
| | |
| == Data Processing ==
| |
| | |
| The primary data reduction package for the EVLA will be the CASA (Common Astronomy Software Applications) package, which will also be used for ALMA. The first public release of the package is now available; see http://casa.nrao.edu for more information.
| |
| | |
| NRAO will continue to support AIPS (Astronomical Image Processing System) for the foreseeable future, at least until VLBI functionality has been incorporated into CASA. See http://www.aoc.nrao.edu/aips/ for more details.
| |
| | |
| It will be possible to reduce data obtained through the Open Shared Risk Observing program using either CASA or AIPS.
| |
| | |
| == Travel Support for Visiting the DSOC and EVLA ==
| |
| | |
| For each observing program scheduled on an NRAO telescope, reimbursement may be re- quested for one of the investigators from a U.S. institution to travel to the NRAO to observe, and for one U.S.-based investigator to travel to the NRAO to reduce data. Reimbursement may be requested for a second U.S.-based investigator to either observe or reduce data provided the second investigator is a student, graduate or undergraduate. In addition, the NRAO will, in some cases, provide travel support to the Observatory for research on archival data. The reimbursement will be for the actual cost of economy airfare, up to a limit of $1000, originating from within the U.S. including its territories and Puerto Rico. Costs of lodging in NRAO facilities can be waived on request in advance and with the approval of the relevant site director. No reimbursement will be made for ground transportation or meals.
| |
| | |
| To qualify, the U.S. investigator must not be employed at a Federally Funded Research and Development Center (FFRDC) or its sponsoring agency. The NSF maintains a master government list of some FFRDCs at http://www.nsf.gov/statistics/ffrdc/.
| |
| | |
| To claim this reimbursement, obtain an expense voucher from Lori Appel in the Assistant Director’s office in the DSOC.
| |
| | |
| == Student Assistance for Data Reduction Visits to the DSOC ==
| |
| | |
| Students visiting the DSOC for the purposes of working on an EVLA or VLBA observing program may be eligible to have their lodging expenses in the NRAO guest house covered by NRAO. To qualify, the student must be a graduate or undergraduate enrolled at a University in the U.S., working on an approved observing program. These are the same qualifications as required for NRAO support of air travel costs described above. In addition, the duration of the visit should be between 5 and 30 days. Requests for support should be made to Claire Chandler at least 4 weeks in advance of the proposed visit. If this is a first time visit then the student should be accompanied by a collaborator on the project, or alternatively an NRAO collaborator may be requested.
| |
| | |
| == Computing at the DSOC ==
| |
| | |
| A primary goal of the computing environment at the DSOC is to allow every user full access to a workstation during his/her visit. There are 10 public workstations available at the DSOC for full-time data reduction by visitors. They are mostly four-processor, 2.0-GHz PCs running Linux.
| |
| | |
| For hardcopy, we have a number of high volume B&W laser printers, two color Postscript laser printers which can reproduce on both paper and transparencies, and one wide-bed color printer.
| |
| | |
| Visitors should reserve time on these workstations when they make their travel arrangements at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml (see also Section 5.12). Note that users may request remote access to the visitor machines as well, without actually visiting the DSOC. Please contact the computing helpdesk (e-mail to helpdesk@aoc.nrao.edu, extension 7213, office 262) for further information about this and any other computing assistance while at the DSOC.
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| | |
| For a more complete description of computing facilities at the DSOC, see http://www.aoc.nrao.edu/computing/.
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| | |
| == Reservations for the EVLA site and/or DSOC ==
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| | |
| Accommodation for visitors is no longer available at the EVLA site. Observers wishing to be present for their observations should stay in the nearby towns of Magdalena or Datil.
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| | |
| Reservations must be made at least 1 week prior to your visit to the NRAO/NM, and 2 weeks’ notice is preferred, through the online form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml. Computing re- quirements and the level of staff assistance needed must be specified through the online form.
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| | |
| First-time visiting students will be allowed to come to the NRAO/NM for observations or data reduction only if they are accompanied by their faculty advisor, an experienced collaborator, or if they have an NRAO staff collaborator.
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| | |
| == Staying in Socorro ==
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| | |
| Visitors to Socorro can take advantage of the NRAO Guest House. This facility contains eight single, four double, and two two-bedroom apartments, plus a lounge/kitchen, and full laundry facilities. The Guest House is located on the New Mexico Tech (NMIMT) campus, a short walk from the DSOC. Reservations are made through the online registration form at http://science.nrao.edu/evla/visitorinfo/visitorregform.shtml.
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| | |
| == Help for Visitors to the EVLA and DSOC ==
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| | |
| We encourage observers to come to Socorro to calibrate and image their data. This is the best way to ensure the quickest turnaround and the best results from their ob- serving. While in Socorro, each observer will interact with members of the DSOC staff in accordance with his/her level of experience and the complexity of the observ- ing program. If requested through the reservation form, the visiting observer will be guided through the steps of data calibration and imaging by a pre-arranged staff friend or scientific collaborator. A list of staff scientists and their interests can be found at http://www.aoc.nrao.edu/epo/ad/aoc-research.shtml. The data analysts, computing helpdesk, and other staff are also available for consultation on AIPS and CASA procedures, and systems questions.
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| | |
| == On-Line Information about the NRAO and the EVLA ==
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| | |
| NRAO-wide information is available on the World Wide Web through your favorite Web browser at URL http://www.nrao.edu, and information specific to astronomers using the EVLA may be found at http://science.nrao.edu/evla/. Information about the VLA may be found at http://www.vla.nrao.edu/astro/. We strongly recommend usage of this on-line service, which is regularly updated by the NRAO staff.
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| | |
| = Miscellaneous =
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| == Publication Guidelines ==
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| | |
| === Acknowledgement to NRAO ===
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| | |
| Any papers using observational material taken with NRAO instruments (EVLA or other- wise) or papers where a significant portion of the work was done at NRAO, should include the following acknowledgement to NRAO and NSF:
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| | |
| :''The National Radio Astronomy Observatory is a facility of the National Sci- ence Foundation operated under cooperative agreement by Associated Universities, Inc.''
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| === Dissertations ===
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| | |
| Students whose dissertations include observations made with NRAO instruments are ex- pected to provide copies of, or links to, their theses for inclusion and maintenance at the NRAO library. These will be catalogued and made available via the NRAO library catalogue.
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| === Preprints ===
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| | |
| NRAO requests that you submit the astro-ph link or an electronic copy of any accepted papers that include observations taken with any NRAO instrument or have NRAO au- thor(s) to the Observatory Librarian. For further information, contact the Librarian in Charlottesville (library@nrao.edu).
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| === Reprints ===
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| | |
| Given the prevalence of electronic access to the literature, NRAO will no longer pay for the purchase of reprints of published papers.
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| | |
| === Page Charge Support ===
| |
| | |
| The following URL contains complete information on the observatory’s policy regarding page charge support: http://www.nrao.edu/library/pagecharges.shtml. The following is a summary:
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| | |
| :• When requested, NRAO will pay the larger of the following:
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| ::– 100% of the page charge share for authors at a U.S. scientific or educational institute reporting original results made with NRAO instrument(s).
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| ::– 100% of the page charge share for NRAO staff members.
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| :• Page charge support is provided for publication of color figures.
| |
| :• To receive page charge support, authors must comply with all of the following requirements:
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| ::– Include the NRAO footnote in the text (see Section 6.1.1).
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| ::– Send the astro-ph link or an electronic copy of the paper upon acceptance or posting on astro-ph to the Observatory Librarian (library@nrao.edu), with the request for page charge support. The Librarian will respond with the amount covered (based on the NRAO page charge policy) and will request the page charge form, with manuscript information completed, via fax (434-296-0278) or e-mail (library@nrao.edu). For questions, contact the Observatory Librarian at 434-296-0254.
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| = Documentation =
| |
| | |
| Documentation for EVLA data reduction, image making, observing preparation, etc., can be found in various manuals. Most current manuals are available on-line (see Section 5.15). 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 EVLA user, not listed here, are available from our website.
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| :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.
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| :2. PROCEEDINGSFROMTHE1998SYNTHESISIMAGINGWORKSHOP:Thisisan 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.
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| :3. INTRODUCTION TO THE NRAO VERY LARGE ARRAY (Green Book): This manual has general introductory information on the VLA. Topics include theory of interferometry, hardware descriptions, observing preparation, data reduction, image making and display. Major sections of this 1983 manual are now out of date, but it nevertheless remains the best source of information on much of the VLA. Copies of this are found at the VLA and in the DSOC, but no new copies are available. Much of this document is available for downloading through the NRAO’s website. It does not include any information about EVLA-specific hardware and software.
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| | |
| :4. *A SHORT GUIDE FOR VLA SPECTRAL LINE OBSERVERS: This describes spec- tral line observations with the VLA. It may be useful for users of the VLA archive for understanding how archival observations may have been set up (local oscillator settings, etc.). It is VLA-specific, however, and much of it does not apply to the EVLA.
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| :5. *AIPS COOKBOOK: The “Cookbook” description for calibration and imag- ing under the AIPS system can be found near all public workstations in the DSOC. The latest version has expanded descriptions of data calibration imag- ing, cleaning, self- calibration, spectral line reduction, and VLBI reductions. See http://www.aoc.nrao.edu/aips/cook.html.
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| :6. *GOING AIPS: This is a two-volume programmers manual for those wish- ing to write programs under AIPS. It is now somewhat out of date. See http://www.aoc.nrao.edu/aips/aipsdoc.html#GOAIPS.
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| :7. *VLA CALIBRATOR MANUAL: This manual contains the list of VLA Calibra- tors in both 1950 and J2000 epoch and a discussion of gain and phase calibra- tion, and polarization calibration. It will remain useful for the EVLA. See http://www.vla.nrao.edu/astro/calib/manual/.
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| :8. *The Very Large Array: Design and Performance of a Modern Synthesis Radio Tele- scope, Napier, Thompson, and Ekers, Proc. of IEEE, 71, 295, 1983.
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| :9. *HIGH FREQUENCY OBSERVING GUIDE. A web-based manual with a great deal of information for users about observing with the VLA 0.7 cm and 1.3 cm observing systems. Seehttp://www.vla.nrao.edu/astro/guides/highfreq/. VLA-specific, not updated for the EVLA.
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| :10. *VLA MEMO SERIES. See http://www.vla.nrao.edu/memos/.
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| :11. *EVLA MEMO SERIES. See http://www.aoc.nrao.edu/evla/memolist.shtml.
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| | |
| :12. *The VLA Expansion Project: Construction Project Book. The EVLA Project Books contains the technical details of the EVLA construction project. It is available online at http://www.aoc.nrao.edu/evla/pbook.shtml.
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| | |
| :13. *CASA COOKBOOK: The CASA “Cookbook” for use of the pack- age for data reduction is available, along with other documenta- tion, from the CASA home page (http://casa.nrao.edu). See http://casa.nrao.edu/Doc/Cookbook/casa cookbook.pdf.
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| | |
| = Key Personnel =
| |
| | |
| {| border="1"
| |
| |+ '''Table 14: Key EVLA Personnel'''
| |
| !Name
| |
| !Phone
| |
| !Room
| |
| !Notes
| |
| |-
| |
| | Lori Appel || 7310 || 340 || Proposal handling administrator, AD office
| |
| |-
| |
| | John Benson || 7399 || 366 || Data archive
| |
| |-
| |
| | Sanjay Bhatnagar || 7376 || 309 || CASA; Imaging algorithms
| |
| |-
| |
| | Walter Brisken || 7133 || 373 || Pulsars; EVLA calibration
| |
| |-
| |
| | Bryan Butler || 7261 || 344 || EVLA Computing Division Head; Planets
| |
| |-
| |
| | Claire Chandler || 7365 || 328 || Deputy AD for Science; EVLA high frequencies
| |
| |-
| |
| | Barry Clark || 7268 || 308 || EVLA/VLBA scheduling; EVLA software
| |
| |-
| |
| | Mark Claussen || 7284 || 268 || EVLA user support; EVLA/VLBA scheduling
| |
| |-
| |
| | Vivek Dhawan || 7378 || 310 || EVLA commissioning
| |
| |-
| |
| | Bob Dickman || 7300 || 336 || EVLA/VLBA Assistant Director
| |
| |-
| |
| | Dale Frail || 7338 || 332 || Pulsars; transient sources
| |
| |-
| |
| | Miller Goss || 7267 || 332 || Spectral line
| |
| |-
| |
| | Eric Greisen || 7236 || 318 || AIPS
| |
| |-
| |
| | Helpdesk || 7213 || 262 || Computer Helpdesk
| |
| |-
| |
| | Leonid Kogan || 7383 || 312 || AIPS
| |
| |-
| |
| | Mark McKinnon || 7273 || 326 || EVLA Project Manager
| |
| |-
| |
| | Joseph McMullin || 7315 || 330 || EVLA Commissioning
| |
| |-
| |
| | Dan Mertely || 7128 || VLA-128 || RFI monitoring and mitigation
| |
| |-
| |
| | Amy Mioduszewski || 7263 || 208 || AIPS; EVLA calibrator models; VLBI at the EVLA
| |
| |-
| |
| | George Moellenbrock || 7406 || 368 || CASA; polarimetry
| |
| |-
| |
| | Emmanuel Momjian || 7452 || 301 || EVLA commissioning
| |
| |-
| |
| | Steve Myers || 7294 || 376 || EVLA calibration; polarimetry
| |
| |-
| |
| | Juergen Ott || 7174 || 369 || Spectral line; Mosaicing; EVLA commissioning
| |
| |-
| |
| | Frazer Owen || 7304 || 320 || High dynamic range; wide-field imaging
| |
| |-
| |
| | Peggy Perley || 7214 || 282 || Deputy Assistant Director for Operations
| |
| |-
| |
| | Rick Perley || 7312 || 362 || EVLA Project Scientist; calibration; polarimetry
| |
| |-
| |
| | Receptionist || 7300 || Front || DSOC receptionist; aocrecep@nrao.edu
| |
| |-
| |
| | Reservationist || 7357 || 218 || DSOC reservations; nmreserv@nrao.edu
| |
| |-
| |
| | James Robnett || 7226 || 258 || Computing Infrastructure Division Head
| |
| |-
| |
| | Michael Rupen || 7248 || 206 || EVLA scientific software; spectral line
| |
| |-
| |
| | Lorant Sjouwerman || 7332 || 367 || EVLA Pipeline; OPT; archives
| |
| |-
| |
| | Ken Sowinski || 7299 || 375 || On-line systems
| |
| |-
| |
| | Meri Stanley || 7238 || 204 || Lead data analyst
| |
| |-
| |
| | Gustaaf van Moorsel || 7396 || 348 || User Support
| |
| |-
| |
| | VLA Operator || 7180 || VLA || On-duty EVLA Operator
| |
| |-
| |
| | VLBA Operator || 7251 || 269 || On-duty VLBA Operator
| |
| |-
| |
| | Joan Wrobel || 7392 || 302 || EVLA/VLBA scheduling
| |
| |-
| |
| | Wes Young || 7337 || 378 || CASA system adminstration
| |
| |}
| |
| | |
| Note: You may e-mail any of the above individuals by addressing your message to “first initial last name”@nrao.edu. Thus, you may contact Joanne Astronomer at: “jastrono@nrao.edu”. The name is trun- cated to eight characters. For contact with the AIPS software group, please e-mail “daip@nrao.edu”. 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 DSOC. If you are calling from the U.S. or Canada, the 4-digit numbers must be preceded with: 1-575-835 from landlines, or 575-835 with mobile phones. If you are calling from outside the U.S. or Canada, the 4-digit numbers must be preceded with 1-575-835.
| |
| | |
| = Acknowledgements =
| |
| | |
| This EVLA Observational Status Summary is based substantially on its predecessor, the VLA Observational Status Summary. Over the VLA history of almost 30 years, many individuals contributed to that document by writing sections, editing previous versions, commenting on draft material, and implementing the capabilities described herein. We thank all these contributors for their efforts. Please contact the listed editors of the present document with questions on the material, or suggestions that would enhance the clarity of this guide.
| |
