# vla antenna positions

You should spend several minutes displaying the data in various formats. Once the download is complete, unzip and unpack the file (within a working directory, where you will later run CASA): There are a number of possible ways to run CASA, described in more detail in Getting Started in CASA. You can inspect this with plotms as we did above. A final example is shown in Figure 3C. With reference to the original self-calibration equation above, if the observed visibility data cannot be modeled well by this equation, no amount of self-calibration will help. Observing with a radio interferometer in the presence of a thick ionosphere is … You can also find the flux density values in the CASA logger: Again, the VLA calibrator manual may be used to check whether the derived flux densities look sensible. The data were acquired with two subbands (spectral windows) around 4.6 and 7.5 GHz. Non-interactively via a script. At any stage in the cleaning, you can adjust the number of iterations that tclean will do before returning to the GUI. When you are happy that all data (particularly on your target source) look good, you may proceed. This will eventually be done using the task fluxscale. The initial display in the logger will include. Antennas other than ea05 look OK. We will not be able to transfer calibration for antenna ea05 so we flag it from the data: For the following bandpass solution we need only solve for our bandpass calibrator, and we will do so now after flagging. You can find the data in the CASA repository. From this display (see Figure 4), you see immediately that the flagging we did earlier of antennas 10 and 12 (ea13 and ea15) has taken affect. the previous clean, we found: and so the peak flux density is 0.157 Jy/beam. Bookkeeping is important! operated under cooperative agreement by Associated Universities, Inc. Alternatively, you can trace out a more custom shape to better enclose the irregular outline of the supernova remnant. This task returns a Python dictionary which we capture in the variable mystat. Task plotms allows one to select and view the data in many ways. We use the CASA task gaincal to solve for phase versus time for the central channels on our three calibrators: To really see what is going on, we use plotms to inspect the solutions from gaincal for a single antenna at a time, iterating through each antenna in sequence by clicking on the Next button (rightward pointing single green arrow) on the GUI to advance the displayed antenna. The procedure is to assume that the flux density of a primary calibrator source is known and, by comparison with the observed data for that calibrator, determine the $g_i$ values (the antenna gains). We are providing this starting data set, rather than the true initial data set for at least two reasons. Interactively examining task inputs. You can quickly see that the last source observed (J0319+4130, a polarization calibrator, shown in purple) is the brightest source in this observation. This supernova remnant has lots of structure - try both standard and multi-scale clean. From this we will make an I-only multiscale image (3c391_ctm_spw0_I.image) -- and in particular the model (3c391_ctm_spw0_I.model) -- to generate a series of gain corrections that will be stored in 3C391_ctm_mosaic_spw0.selfcal1. numbers given in listobs. The importance of this step is that the visibility function is a function of $u$ and $v$. You should be able to process this data in a very similar manner to the C-band data on 3C391. The most straightforward statistic is the peak brightness, which is determined by imstat. The CORRECTED_DATA column of the MS now contains the self-calibrated visibilities, they will now be used by tclean. Second, while necessary, many of these steps are not fundamental to the calibration and imaging process, which is the focus of this tutorial. There are four basic antenna arrangements, called configurations, whose scales vary by the ratios 1 : 3.28 : 10.8 : 35.5 from smallest to largest. For the purposes of this tutorial we have created a starting data set, upon which several initial processing steps have already been conducted. We will use a multi-scale cleaning algorithm because the supernova remnant contains both diffuse, extended structure on large spatial scales and finer filamentary structure on smaller scales. myset = setjy(...). This approach will do a better job of modeling the image than the classic clean delta function. Both to get a sense of the array, as well as identify an antenna for later use in calibration, use the task plotants (see Figure 1). To create a clean box (a region within which components may be found), hold down the right mouse button and trace out a rectangle around the source, then double-click inside that rectangle to set it as a box. The positions of EOVSA antennas are determined using observations by the 27-m (Ant 14) low-frequency receiver (S and C band) of celestial radio sources during several observation runs in fall 2016. (Angular scale [in radians] ~ 1/baseline [in wavelengths]. These data are from the D configuration, in which antennas are particularly susceptible to being blocked (shadowed) by other antennas in the array, depending upon the elevation of the source. This step solves for the complex bandpass, $B_i$. This document describes the procedure followed and the final? Even though each of the target scans is on the same source (3C391), the observation is done as a mosaic of 7 fields (see the listobs output above). Self-calibration assumes that the model is perfect. The K gain type in gaincal solves for the relative delays of each antenna relative to the reference antenna (parameter refant), so be sure you pick one that is there for this entire scan and good. In this tutorial, we will run the cleaning task interactively so that we can set and modify the mask: Task tclean is powerful with many inputs and a certain amount of experimentation likely is required. You can also employ the region panel to save a region you have created for later use. Work at your own pace, look at the inputs to the tasks to see what other options exist, and read the help files. At this stage in the calibration, we have not yet solved for the flux density scaling. When you are happy with the clean regions, press the green circular arrow button on the far right to continue deconvolution. Even after the initial calibration using the amplitude calibrator and the phase calibrator, there are likely to be residual phase and/or amplitude errors in the data. NRAO monitors the positions of the VLA antennas on a regular basis. The weak thermal emission from the largest minor planets can be detected and measured at all points around their orbits at microwave frequencies using the Very Large Array (VLA). It is always a good idea to examine the data before jumping straight into calibration. In this mode, one types. The VLA antenna contractor was E-Systems, Garland Division, Dallas, Texas. VLA polarization calibration Instrumental. We have chosen to colorize by scan; it's clear that the slopes are steady over time. extended atmosphere with the VLA in A conﬁguration and the VLBA Pie Town antenna to obtain maximum spatial resolution. position of the interferometer elements (the antennas or stations) and the viewing direction. This CASA guide describes the calibration and imaging of a multiple-pointing continuum data set taken with the Karl G. Jansky Very Large Array (VLA) of the supernova remnant You can adjust the color scale and zoom in to a selected region by assigning mouse buttons to the icons immediately above the image (hover over the icons to get a description of what they do). None of these antennas should be chosen as the reference antenna during the calibration process, below. The index '1' above refers to the field number. For the 1994 September observations, St. Croix had water in the 15 GHz feed. The initial Science Data Model (SDM) file was converted into a measurement set. Again, the scatter is normal at this pre-calibration stage. We explore a few possible algorithms for finding four to six antenna positions for the VLA Upgrade's A + configuration, which will bridge the gap between the VLA's A configuration and the inner VLBA antennas. Finally, it is common for the array to require a small amount of time to settle down at the start of a scan. Using this technology, scientists can study astronomical phenomena that would otherwise be invisible. Upper left: Elevation vs Time of the modeled source (blue). Images 19A through 19B show that there is no sign of bad data remaining. We chose the Hot Metal 1 colormap and set the Scaling Power Cycles to -1 to better emphasize the faint emission and compare to the noise (Figure 25). Although CASA has the feature that its Fourier transform engine (FFTW) does not require a strict power of 2 for the number of linear pixels in a given image axis, it is somewhat more efficient if the number of pixels on a side is a composite number divisible by any pair of 2 and 3 and/or 5. In this example, we have elected to show phase as a function of (frequency) channel for a single baseline (antenna='ea01&ea21' ) on the bandpass calibrator. NRAO also provides both formal and informal programs in education and public outreach for teachers, students, the general public, and the media. You may notice that there are antenna-to-antenna variations (under the Display tab select Colorize by Antenna1). Very Large Array ECE researchers working with the National Radio Astronomy Observatory (NRAO) have successfully demonstrated a new antenna feed system for the Very Large Array (VLA). As mentioned above, restarting tclean with different multiscale=[...] choices can help also. The parameters are similar as before. self calibration (see Section 5.11), but we have lectures on Self-calibration given at NRAO community days. Open viewer and use it to display the corrected image (Figure 26). These allow you to make changes to the plotting selection without having to re-launch plotms. Then trace out a shape by right-clicking where you want the corners of that shape. The setting applymode='calflagstrict' will be even more stringent about flagging things without valid calibration, while applymode='calonly' will calibrate those with solutions while passing through data without unchanged. The first option is simplest as it is the same object using a different spectral window; for a more rewarding challenge try the L-band dataset on G93.3+6.9. Most tasks in are designed to work for a generalized interferometer or single-dish. As expected, the bandpass phases are relatively flat (see Figure 8B), with the slopes (Figure 3C) removed by the delay calibration. The recorded data are then sent to Socorro, NM to be processed by a powerful computer known as a correlator Correlator A specialized supercomputer that multiplies the data from two antennas and averages the result over time. Positions were obtained in the ICRF directly through phase referencing of the stars to nearby ICRF quasars whose positions are accurate at the 0.25 mas level. Parameter iteration='antenna' is used to step through separate plots for each antenna. As part of this radio study we have also obtained new multi-epoch radio positions. In the D configuration, the fringe rate is relatively slow and time-average smearing is less of a concern. For most antennas, we see a fairly smooth variation with time, so we expect to be able to calibrate the data nicely. Each antenna lives on a mount next to the rails used to change the antenna positions in the array. Figure 3B shows this example, including time averaging of '1e6' seconds (any large number that encompasses more than a full scan will do). As another example of using plotms for a quick look at your data, select the Data tab and specify field 0 (source J1331+3030, a.k.a. V'_{ij} = G_i G^*_j V_{ij} Abstract: We have used the Very Large Array (VLA), linked with the Pie Town Very Long Baseline Array antenna, to determine astrometric positions of 46 radio stars in the International Celestial Reference Frame (ICRF). Antenna 3 (ea04) is missing the last scan and antenna 23 (ea26) is missing scans near the end. In the logger you can see the corrections reported: This particular set of observations was taken 24 April 2010, so the corrections shown above are for antennas that were moved BEFORE that date, but whose updated positions were not placed into the online database until later. You can also set multiple clean regions. This lets astronomers place them … As discussed above, the absolute magnitude of the gain amplitudes ($g_i$) are determined by reference to a standard flux density calibrator. In brief, there are at least three different ways to use CASA: If you are a relative novice or just new to CASA, it is strongly recommended to work through this tutorial by cutting and pasting the task function calls provided below after you have read all the associated explanations. You will need to find it (e.g., using, We have not edited out bad or dead antennas for you (unlike in 3C391). There is a small amount of discussion in the old CASA Reference Manual on However, assuming that the calibrator was sufficiently close to the target, and the weather was sufficiently well-behaved, then this is a reasonable approximation and should get us a sufficiently good calibration that we can later use self-calibration to correct for the small inaccuracies thus introduced. We should now have fully-calibrated visibilities in the CORRECTED_DATA column of the measurement set, and it is worthwhile pausing to inspect them to ensure that the calibration did what we expected it to. This tutorial is made up of such calls, which were developed by looking at the inputs for each task and deciding what needed to be changed from default values. combine the two calibrated MSs in tclean to make a deeper MFS image (this might be tricky). The ratio of amplitude solutions between the two sources will be used in a later calibration step (fluxscale) to determine the actual flux density of J1822-0938. The antennas can be referenced using either convention; antenna='22' would correspond to ea25, whereas antenna='ea22' would correspond to ea22. ), By contrast, if you make a similar plot for field 8 (one of the 3C 391 fields), the result is a visibility function that falls rapidly with increasing baseline length. This is also a good time to check that our chosen reference antenna (ea21) has good phase stability (i.e., the phase difference between the right and left polarizations is stable with time). You will have to avoid it through channel ranges and/or edit it out. Before starting the calibration process, we want to get some basic information about the data set. This term is used where the CW and CCW wraps overlap between 275° and 85° real AZ, passing through 0° / 360° (north) (CW: 275° to 445° (85° real AZ) / CCW: −85° (275° real AZ) to 85°). From this, we find the following: Before beginning our data reduction, we must start CASA. This allows us to find the true flux density of the secondary calibrator. The data were averaged from the initial 1-second correlator sample time to 10-second samples. We will run this task here on the newly calibrated and split-out data set before moving on to imaging. Below, the current VLA configuration is listed with a visual representation of that particular configuration. The first command below shows the amplitude solutions (one per polarization) and the second command below shows the phase solutions (one per each polarization). North. If you find you don't like your region you can dismiss it with with ESC key or using the remove region "X" button in lower right of the panel. The supernova remnant itself is known to have a diameter of order 9 arcminutes, corresponding to about 216 pixels for the chosen cell size. This will include the flux density value within the region selected. This is known as 'quack' flagging. As mentioned above, we can guide the clean process by allowing it to find clean components only within a user-specified region. The free electron content of the ionosphere varies with time of day, season, geographic latitude and Solar activity. The most effective cleaning occurs with at least 4-5 pixels across the synthesized beam. How big were the phase changes made by the calibration? This expression also neglects other factors, such as the shape of the primary beam. To examine the observing conditions during the observing run, and to find out any known problems with the data, download the observer log. Each station consists of a 25-meter radio antenna dish and a control building. Please contact the NRAO Helpdesk. In general, for calibration purposes, one would like to select an antenna that is close to the center of the array (and that is not listed in the operator's log as having had problems!). For example, plot (with colorization by polarization) for the first block of 3C286 data only: The first stage of bandpass calibration involves solving for the antenna-based delays which put a phase ramp versus frequency channel in each spectral window (Figure 3C). Carpentry. As you clean you will see faint sources all over the field; welcome to L-band imaging. If the data set has more than one spectral window, depending upon where they are spaced and the spectrum of the source, it is possible to find quite different flux densities and spectral indexes for the secondary calibrators. (See the CASAguide on radio galaxy 3C75 for an introduction to polarization calibration.). Solving for the bandpass won't hurt for continuum data, and, for moderate or high dynamic range image, it is essential. Finally, we apply the calibration to the target fields in the mosaic, linearly interpolating the gain solutions from the secondary calibrator, J1822-0938. Any updated positions that were entered into the database BEFORE your observations will already be accounted for in your data. The tclean task naturally operates in a flat noise image, i.e., an image where the effective weighting across the mosaic field of view is set so that the noise is constant. For this analysis, it is better to use the version of the viewer that is run from the OS command line rather than the CASA command line. The Double click inside of that region (using the same mouse button used to make the region), and the statistics will be reported. In general, if the noise is well-behaved in the image, when averaged over a reasonable solid angle, the noise contribution should approach 0 Jy. We make some standard plots (see Figures 19A through 19D): Inspecting the data at this stage may well show up previously-unnoticed bad data. Appendix V VLA Maintenance Tasks in . Antenna ea10 was not in the array, but, for the other two antennas, any improved baseline positions need to be incorporated. The flux density in each Stokes plane (IQUV) for the reference channel 0 is reported, followed by the I flux density in each of the channels to be used for scaling the data. Later, for the final step in determining the calibration solutions, we will use the calibrated gains of the two calibrator sources to transfer the flux density scaling to the secondary gain calibrator (J1822-0938). The baseline length at which the visibility function falls to some fiducial value (e.g., 1/2 of its peak value) gives a rough estimate of the angular scale of the source. 34-m antenna at diﬀerent elevation positions. Task fluxscale will print to the CASA logger the derived flux densities of all calibrator sources specified with the transfer parameter. In order for the amplitude of 3C 286 in the data to match the amplitude of its model (which we set above with setjy), little scaling of the solution is required (value = 1.0). In our case, you will find the return value in the CASA command line window: If desired, this can be captured by calling the task by setting it to a variable, e.g. Example output would be. The complex gain calibrator (J1822-0938, shown in magenta) is slightly brighter than the target fields. ': One can choose the function assigned to each mouse button; after zooming into the desired view, assign polygon region to a desired mouse button (e.g., left button) by selecting the polygon tool to create the polygonal region as shown in Figure 26 with the desired mouse button. Is this better than the original multiscale image? One final useful plot we will make is a datastream plot of the antenna2 in a baseline for the data versus ea01, using spw 0 and channel 31. Some of the tracks cross the New Mexico highways, as seen during this move in the 1970s. The observations were taken with a full-polarization correlator setup and include a polarization calibrator. This means that the position of the EVLA antennas (and VLA antennas, for that matter) in your data may not be the best available one, which can lead to phase errors. V.4 Improving antenna surfaces: holography. In our example we find a total Flux density of 9.4 Jy. Here, for generality, we show the visibility as a function of frequency $f$ and spatial wave numbers $u$ and $v$. In “holography” data from the EVLA, reference antennas are pointed towards the nominal source position, as normal.All other antennas are pointed to an offset position controlled by the observer. One task to examine the data themselves is plotms. calibrated antenna positions. After about 14000 iterations (Figure 24) the residuals were looking good (similar noise level inside and outside of the clean region). After many cycles, when only noise is left, you can hit the red-and-white stop-sign icon to stop cleaning. At each epoch, nine of the 10 VLBA antennas were operating at 15 GHz. Here, we have put together a collection of exclusive video tours we call the VLA Explorer. A not-uncommon limitation for moderately high dynamic range imaging is that there may be baseline-based factors that modify the true visibility. The wheels are mounted on assemblies at each corner of the vehicle that can rotate to allow "turning" the 90-degree rail intersections that connect each antenna mounting station with the main rail line for each arm of the VLA's "Y" layout.. Automotive. To get a more detailed view of the central regions containing the emission, zoom in by first left clicking on the zoom button (leftmost button in third row) and tracing out a rectangle with the left mouse button and double-clicking inside the zoom box you just made. When enough precise positional observations have been accumulated, the orbits of the … The wheels are mounted on assemblies at each corner of the vehicle that can rotate to allow "turning" the 90-degree rail intersections that connect each antenna mounting station with the main rail line for each arm of the VLA's "Y" layout.. First, many of these initial processing steps can be rather time consuming (> 1 hr). has 27 active antenna dishes. By contrast, for the rms noise level, one can load the original (un-pbcor) image: and to exclude the source's emission to the extent possible as shown in Figure 27, as the source's emission will bias the estimated noise level high. The data are color coded by spectral window, as earlier plots from plotms indicated that the visibility data are nearly constant with baseline length. "2@FD" for the Fort Davis, Texas, antenna. The $u$ and $v$ coordinates are the baselines measured in units of the observing wavelength, while the $l$ and $m$ coordinates are the direction cosines on the sky. Therefore, we use impbcor to divide the .image by the .pb image to produce a primary beam corrected restored image: You can open this in the viewer and see that it has indeed raised the noise (and signal) at the edges of the mosaic. This may or may not be a good thing. In the example above, the descriptor, A common metric for self-calibration is whether the image. Be careful when making images and setting clean regions or masks. You want to be sure to capture as much of the source total flux density as possible, but not include low-level questionable features or sub-structure (ripples) that might be due to calibration or clean artifacts. To apply the calibration we have derived, we specify the appropriate calibration tables, which are then applied to the DATA column, with the results being written in the CORRECTED_DATA column. For the pre-upgrade VLA, the ultimate flux density scale at most frequencies was set long ago by observations of 3C 295. To do this, we iterate over the different calibrators, in each case specifying the source to be calibrated (using the field parameter). To motivate the need for solving for the bandpass, consider Figure 7. Self-calibration requires experimentation. Task tclean will make several output files, all named with the prefix given as imagename. Data calibration will take care of much of that scatter. You can open this from inside CASA using '! If updated positions were entered into the database AFTER your observation date, the corrections to the newly measured positions can still be applied during your data reduction process in this step. In the first step, we derive the appropriate complex gains $g_i$ and $\theta_i$ for the flux density calibrator 3C 286. The data were taken in early science shared-risk observing mode, with 128 MHz of bandwidth in each of two widely spaced spectral windows, centered at 4.6 and 7.5 GHz. Drafting. To do this, we use the task fluxscale, which produces a new calibration table containing properly-scaled amplitude gains for the secondary calibrator. Will it be possible to remember one month later (or maybe even one week later!) This page was last edited on 29 October 2020, at 14:58. The beam characteristics are repeatable from antenna to antenna. image) are Fourier transform pairs. The next step is to derive corrections for the complex antenna gains, $g_i$ and $\theta_i$. The task listobs can be used to get a listing of the individual scans (set amounts of time devoted to specific targets) comprising the observation, the frequency setup, source list, and antenna locations. In spite of RFI, the antenna-based calibration is remarkably resilient to moderate-to-low RFI contamination (which tends to be baseline-based). Such a visibility function indicates a highly resolved source. The variations of the ﬁrst two natural frequencies with respect to the antenna elevation position are shown in Fig. Stepping through to that antenna reveals Figure 5. The Basics. gaincal step will report a number of solutions with insufficient SNR. It is important to have an idea of what values to use for the image pixel (cell) size and the overall size of the image. To do that, right-click on the closed polygonal icon. [8] Because of the consequences of the Y-shape of the VLA… This shows, assuming that ea01 is in the entire observation, when various antennas drop out. The reduction strategy is to determine various corrections from the calibrators, then apply these correction factors to the science data. However, for some kinds of polarization calibration or in extreme situations (e.g., there are problems with the scan on the flux density calibrator), it can be useful (or necessary) to set the flux density of the source explicitly. which indicates that, among other things, the flags that existed in the data set prior to this run will be saved to another file called flagdata_1. More specifically, the observed visibility data on the $i$-$j$ baseline can be modeled as, [math] Antenna Mechanics. More often during reconfigurations then zoom and pan to explore the array by default, we focus... Last edited on 29 October 2020, at 14:58 Garland Division, Dallas, Texas 4 Hobbs. Having to re-launch plotms. ) basis -- roughly monthly, and can be downloaded as a 7-pointing mosaic the. 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