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The new EVLA correlator is extremely powerful in its spectral capabilities. Up to 4 million channels can be observed with a spectral resolution in the Hz regime. Here is a guide to access that spectral line power.

Observation Planning

The wide bands of the EVLA allow users to observe up to 8GHz spectral bandwidth at a time. This will result in extreme continuum sensitivity. In addition, it opens up the possibility to observe one or more spectral lines at a given time. So the first to is to carefully plan the observations.

There are a number of online tools available that help spectral line observers. To start with, the rest frequency of the sopectral line needs to be determined. This can be done with Splatalogue which contains data from the Lovas catalog, the JPL/NASA molecular database], the Cologne Database for Molecular Spectroscopy and others. Note that in addition to molecular line transitions, splatalogue also contains radio recombination lines.

Line Frequency

The first step is to determine the observing frequency $\nu$ of the spectral line. This is derived from the radial velocity $v$ of the source and the rest frequency $\nu _{0}$ of the spectral line.

A full relativistic calculation shows that the velocity $v$ is determined via $v={\frac {\nu _{0}^{2}-\nu ^{2}}{\nu _{0}^{2}+\nu ^{2}}}$. This is called the relativistic velocity. This equation is a bit cumbersome to use and in astronomy to different approximations are typically used instead:

• Optical Velocity $v^{optical}={\frac {\lambda _{0}-\lambda }{\lambda _{0}}}\,\,c=cz$ ($z$ is the redshift of the source)

• Radio Velocity $v^{radio}={\frac {\nu -\nu _{0}}{\nu _{0}}}\,\,c={\frac {\lambda _{0}-\lambda }{\lambda }}\,\,c\neq v^{optical}$

The radio and optical velocities are not identical. At low velocities that difference is small but $v^{optical}$ and $v^{radio}$ diverge more and more for large values.

At significant redshifts, it is possible to place the zero point of the velocity frame into the source via

$\nu =\nu _{0}/(z+1)$

The redshifted $\nu$ can now be used as the input rest frequency for the observations. The velocites that are derived based on such a redshifted rest frequency will then also be correctly scaled for the spread of the velocity scale that is due to the redshift.

Velocity Frames

There are various rest frames which might be appropriate. The following table lists their name, the motion for which one has to correct in order to reduce an observed velocity to that particular rest frame, and the magnitude of the velocity correction. Each subsequent rest frame is obtained by adding the effects of the preceeding ones: Correct for: Amplitude (km s$^{-1}$) Rest frame: Nothing added 0.0 topocentric

Earth's rotation $\leq$0.5

Earth's motion around $\leq$0.013 km s$^{-1}$ geocentric earth/moon barycenter

Earth's motion around $\leq$30 km s$^{-1}$ heliocentric(z) the Sun

Solar motion around the $\leq$0.012 km s$^{-1}$ barycentric Solar System barycenter

Solar motion $\sim$20 km s$^{-1}$ local standard of rest (LSR)

Galactic rotation $\sim$300 km s$^{-1}$ galactocentric

The most commonly used rest frames are heliocentric (to be precise, barycentric is used at the VLA) and local standard of rest (LSR). LSR is generally used in Galactic astronomy and heliocentric in extragalactic astronomy, although the latter is often reduced to galactocentric.

More often than not, one wishes to specify the velocity of the object and let the on-line system do the conversion to sky frequency for each scan. This is called ``Doppler tracking. The user will have to specify, in observe, the rest frequency, choice of rest frame, and radial velocity. Doppler tracking is not implemented during a scan as the frequencies are set at the beginning of each scan. If very accurate tracking is required, one is advised to use short scans. Note that the on-line system uses the same algorithm as dopset whereas observe (version 3 and higher) uses a slightly different method and calculates the observing frequency to within a few tens of meters per sec to the values derived using dopset.

Ideally, one wants calibrators observed at the same sky frequency as the sources. This can be achieved by specifying ``no change in observe for the flukesynthesizer on the calibrators instead of a velocity. The effect of this is that the LO settings are not changed from what they were during the previous scan. If one wants to start a sequence with a calibrator, it is necessary to precede it with a dummy 1-minute source scan to force the on-line computers to set the LO chain to the proper values. WARNING: if the system crashes and comes back up in the middle of a calibrator scan this scan will be useless because the frequency setting will be in error. The VLA Operator should be alerted to the use of the ``no change option, so he can restart an interrupted observation by including again a 1-minute dummy scan. Because of this overhead, it is advised that the use of ``no change be restricted to those cases in which it is essential, such as for obtaining high spectral dynamic range; in general, the frequency difference for nearby calibrators is negligable.

Setting up Correlator Modes

The OPT

spectral line setup

- planning - PST - OPT - exposure calculator

- OPT preparation and setup

 . suckouts
 . filters
 . tuneability of subbands
 . channels full/dual/single pol
 . baselineboard stacking
 . recirculation
 . gapfree setups
 . data rate considerations
 . setups with variable bandwidths

- bandpass

- post processing - hanning smoothing - velocity systems - continuum subtraction

Spectral Line Observing

Current->Revised OSS Guidelines

• An Overview of the EVLA

(last paragraph) The EVLA correlator will be extremely powerful and flexible. Details of the correlator configurations being offered for EVLA early science during the period Sep 2011 - Dec 012 (a full D→A configuration cycle) are described in Correlator Configurations. 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 Correlator Configurations.

• Limitations on Imaging Performance

(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.

• Correlator Configurations

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.

Detailed Guidelines

Observing Preparation Recommendations

Calibration Strategy

• Bandpass Setup

All observations with the EVLA---even those with the goal of observing continuum---require bandpass calibration. A bandpass calibrator should be bright enough, or observed long enough, so that the bandpass calibration does not significantly contribute to the noise in the image. This implies that, for a bandpass calibrator with flux density S_cal observed for a time t_cal and a science target with flux density S_obj observed for a time t_obj, $S_{cal}{\sqrt {t_{cal}}}$ should be greater than $S_{obj}{\sqrt {t_{obj}}}$. How many times greater will be determined by one's science goals and the practicalities of the observations, but $S_{cal}{\sqrt {t_{cal}}}$ should be greater by at least a factor of two.

The bandpass calibrator should also be a point source or have a well-known model. At low frequencies, the absolute flux density calibrators (3C48, 3C147, or 3C286) are quite bright and in many cases can double as the bandpass calibrator. However, at high frequencies, these sources have only moderate flux densities of ~0.5--3 Jy, translating into a potentially noisy bandpass solution.

The stability of bandpasses as a function of time is of concern for high-dynamic-range spectral work. We have found that most antennas show bandpasses that are stable to a few (~2--4) parts in a thousand over a period of several (~4--8) hours [L BAND?].

Dramatic jumps in the bandpass structure (of order a few parts in a hundred) can occur at attenuator changes. The observer can track down such attenuator changes in their data using the switched power information; the On - Off power ('PDIF' in AIPS) wil show a clear discontinuity. For this reason, it behooves the spectral line observer to observe a bandpass calibrator at least twice during their observations. Multiple observations will provide a check that all is well on most antennas and a mechanism for identifying any "problem" antennas. However, we do not expect that interpolating in time between consecutive bandpass solutions will bear much fruit for the observer. The low-level variations observed on some antennas tend to not be smooth functions of time and will likely not be corrected with interpolation.

If there is only one observation of the bandpass calibrator, the observer should be careful to minimize the number of shadowed antennas, as an antenna without a bandpass determined for it will essentially be flagged for the rest of the observation.