The JCMT-CSO Interferometer

Over the last few years work has been underway to link the JCMT and its neighbour, the 10.4 m Caltech Submillimeter Observatory together to form the first astronomical interferometer operating at submillimetre wavelengths. The telescopes are located at an altitude of 4100 m on the summit of Mauna Kea, Hawaii, separated by a baseline of length 164 m, giving a fringe-spacing of 1.1 arcseconds at 345 GHz.

The progress of this difficult project has been charted in the pages of the Newsletter. Beginning with semester 94B the interferometer was made available to the general user community of observational programmes. Since any use of the system requires the use of both telescopes and the special expertise of its builders, there are some particular conditions attached to applications for time. To bring together all the necessary requirements the interferometer is scheduled in much the same way as the MPE high- frequency receiver (RxG); that is, typically for a period of a couple of weeks per semester. Within this period, successful programmes are scheduled flexibly, depending on weather conditions and equipment availability.

There follows brief information on the technical design, the baseline geometry, the correlator, the system temperature and sensitivity, the effects of the atmosphere on observing, and the data-reduction process. Application procedures for those wishing to make use of the interferometer and complete information on the JCMT/CSO interferometer is also available on the JCMT homepage of the WWW.

Technical outline

The interferometer uses the existing heterodyne receivers at the two telescopes to down-convert signals to an Intermediate Frequency (IF) in the range 1-2 GHz. The two local oscillators must be highly coherent for interferometry to be possible, so they are phase-locked to a common reference frequency that is transmitted from the CSO to the JCMT over a wideband fibre-optic link. Optical fibres are also used to deliver the two IF signals from the receivers to the digital correlator (the DAS) at the JCMT. In general, astronomical signals arrive at one antenna before the other, and delay compensation is thus needed before correlation. Coarse delay compensation is realised by switching appropriate lengths of optical fibre into the signal paths. Finer corrections are made in the software during post-processing of the data. The rotation of the Earth causes the sources to move through the fringe pattern of the interferometer producing a fringe rate of up to ~15 Hz for observations at 345 GHz. Numerically-controlled oscillators are used to generate offset frequencies which are introduced into the LO and IF chains in such a way that this fringe rate is removed. Phase-switching at a fixed rate is also introduced which helps to remove any offsets in the correlator.

This arrangement naturally ensures that only one of the two sidebands generated by the receivers is correlated. The fringe rate can only be removed precisely for one frequency (chosen to be the centre of the IF band) and this sets a limit on the maximum integration time for individual records of around 10 seconds.

The interferometer uses heterodyne receivers in the 220 - 270 GHz and 318 - 360 GHz bands. Observations at around 460 - 500 GHz may also be possible, but the performance in this band is very uncertain and will be extremely weather dependent. Applications requiring this band will be considered but the chances of them being carried out are relatively low. Note that because the interferometer provides its own LO chain, the frequency coverage of the receivers available at JCMT/CSO will likely be somewhat different than is the case for single-antenna work.

Baseline geometry, resolution and applications

The components of the baseline (going from the CSO to the JCMT) are: +43.38 m in the polar direction, - 157.91 m in an east-west direction, and -1.93 m towards the celestial equator. The fringe-spacing is given by l divided by the length of the baseline projected perpendicular to the source. The minimum fringe spacing (1.1" at 345 GHz) therefore occurs when the source passes through the plane perpendicular to the baseline.

With a single baseline, the imaging ability of the interferometer is of course very limited. However the visibility of the fringes as a function of the length and orientation of the projected baseline can be measured and this can provide a good estimate of the size of the source and some information on the shape. By measuring the phase of the fringe, it may eventually be possible to measure positions of objects relative to known sources like quasars, to an accuracy of perhaps 0.1", but this has not yet been achieved. It is however quite easy to measure the relative phases of different spectral line components in a single source, provided they fall within the bandpass of the correlator. These phase differences give information about the angular separations between the different components.

The correlator

The Dutch Autocorrelation Spectrometer (DAS) at the JCMT is used to form the complex cross- correlation function of the two IF signals. There are 2048 lags which produce 1024 complex frequency channels (although only about 800 of these contain good data because of the bandpasses of the filters). The lags can be distributed in many different ways amongst up to 8 sub-bands, each with a bandwidth of about 125 MHz. Some examples of the configurations available are: 800 channels covering 125 MHz with a resolution of 0.156 MHz; 100 channels in each of 8 sub-bands with 1.25 MHz resolution — total bandwidth covered 920 MHz (there is some overlap between the sub-bands); 400 channels covering 500 MHz with a resolution of 1.25 MHz, and the remaining 400 channels covering 125 MHz at 0.313 MHz resolution.

The system temperature and sensitivity

The system temperature is the geometric mean of the JCMT and the CSO system temperatures. These are determined by the standard method by measuring the power from the sky and an ambient load, so that the contribution due to the atmosphere is taken into account. Single-sideband values for operation at 345 GHz are presently in the range 500 K to 1000 K. The value is lower at 230 GHz, but increases going to low elevations and near to the main atmospheric absorption lines. The conversion factor from effective antenna temperature to units of flux is expected to be about 50 Jy/K, assuming an efficiency of 50%. The true value on a given observing run is determined by observing quasars as flux calibrator sources.

(1)

(2)

The sensitivity of the interferometer to continuum emission after integrating for time t(int) over bandwidth dn is given by Equation 1.

For a spectral line of width dv the corresponding result is given by Equation 2.

The interferometer starts to resolve sources of size > 0.5". A thermal source 0.5" x 0.5" must have a brightness temperature of at least 12 K to produce a flux of 260 mJy at 345 GHz.

The integration time of 10 seconds implicit in the above is the shortest which is normally used. For most purposes the records would be added together (after correcting for the effects of the changing delays) to give longer integrations. There is however a limit to this which is set by phase drifts due to uncertainties in the baseline, thermal effects an the atmosphere. This is quite often no longer than 100 seconds. To continue integrating coherently for longer than that it is necessary to have a source in the beam giving a signal to noise ratio of greater than 1 which can be used as a phase reference.

The effects of the atmosphere

The interferometer is more sensitive to the weather than the individual antennae. The Earth's atmosphere introduces random phase fluctuations that decorrelate the signals. The rms. phase fluctuation depends on the weather conditions, the observing frequency and the elevation. There is a strong diurnal effect because the heating of the ground by the sun drives much of the turbulence. In good weather, observations may be possible between ~5pm and 7am, but often the start has to be delayed by an hour or two while the atmosphere settles. Observing beyond 7am is limited because direct sunlight must not be allowed to fall on the CSO surface. A typical night-time rms. Phase fluctuation value is 30 degrees for observations at 345 GHz. Sources can be observed down to elevations of ~25 degrees above the horizon, allowing up to 8 hours of observations on a source of moderate declination. In practice the observations of the main source must be interleaved with those of a quasar for phase, flux and passband calibration.

Data reduction

The output of the DAS is processed using software written specifically for the JCMT-CSO Interferometer. This performs the fine delay correction, the passband correction and has facilities for adding and smoothing spectra. The final output is usually either a complex spectrum or an average flux and phase for the specified part of the passband. The output can be written out in the form of a UVFITS file for further processing with other software packages.

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