Astron. Astrophys. 346, 181-189 (1999)

2. Experimental description

2.1. Experimental context

The observations were carried out with the two 45 cm telescopes of the IOTA interferometer located on Mount Hopkins, in Arizona. A complete updated description of IOTA is given by Traub (Traub 1998). The telescopes are made of a siderostat and a 1/10 afocal beam compressor. After compression, the direction of the beams is corrected for the rapid tip-tilt motion introduced by the atmosphere. This servo looped stabilization is based on the stellar image centroid position measured by visible star trackers, as presented in Fig. 1. One of the beams feeds a first (fixed) long delay line, that compensates for most of the external path difference, and then a second delay line with a shorter stroke and a very smooth motion to enable continuous modulation of the optical path difference during the observations. Thus the interferometer delivers in its central laboratory a pair of stabilized and properly delayed afocal beams which in our case are combined on the TISIS table.

Fig. 1. Optical lay out of the TISIS combination table in the L band. A single X coupler is used for coherent combination of the beams coming from the 2 telescopes. There is only one interferometric output and no monitor of the signals coupled into the fibers.

The telescopes can be relocated over an L shaped set of stations providing baselines ranging in length from 5 to 38 m. Observations were carried out with the South siderostat at its 15m station and the North siderostat at its 15 m station, so that the resulting baseline was 21 m long and quasi aligned with the North-South direction. As far as the interferometer is concerned, the only change with respect to observations in K is the removal of the windows (used to close the delay line and work under vacuum), due to their poor transmission above 3.5 microns. The delay line was then operated under atmospheric pressure.

2.2. Optical layout of the TISIS combination table

The optical scheme for combination in TISIS (Fig. 1) is a simplified version of the FLUOR instrument. Nevertheless a few differences can be noted in the optical arrangements of TISIS:

Fluctuations of the injection efficiency are not monitored at each input fiber, as should be done for optimum accuracy in the visibility computation. A few couplers available from the FLUOR K band experiment were tested and the only one properly balanced with respect to the two input signals in L is an X coupler with a broken output (a "Y"coupler). There is then only one usable -interferometric- output.

There is no polarization control through macrobending of the fibers. In the X-shaped coupler, originally designed and optimized for operation in the K band, the cutoff wavelength of the fundamental guided mode is 1.9 µm. For a radiation with a wavelength typically twice as long, light is not guided as well (Neumann 1988), and important losses occur if the fiber is substantially bent. Consequently, the fibers are just kept as straight as possible to avoid additional losses.

Care has to be taken in order to limit the amount of thermal background received by the detector. A 2 mm cold field stop is then used in order to control the beam etendue, and the corresponding background seen by the detector. A Fabry lens located inside the cold part of the detector is used to reimage the fiber's output on the diaphragm and insure a good filling of the detector's InSb photodiode. With this optical arrangement, the beam etendue seen is defined as the product of the cold aperture area by the acceptance solid angle of the detector (f/130). This is equivalent to 10 at the center of the L band. In comparison, because of the single-mode character of the fibers, the starlight beam etendue seen by the detector is restricted to , which corresponds to the diffraction limited case. The 10 beam etendue seen by the detector already keeps the background fluctuations below dark current noise (Sect. 3.1.1). Reducing it further to would require a still smaller field stop, and very constraining optical alignments in order to keep the starlight on the detector.

We had to use an electrical offset before signal amplification in order to avoid saturation. The overall detection scheme is summarized in Fig. 2.

Fig. 2. Detection scheme

2.3. Observing procedure

Interferograms are obtained as scans around the zero optical path difference. For all the observations the speed of the delay line is set in order to yield an overall fringe speed of 1.5 mm/s, so that the apparent mean fringe frequency is about 400 Hz, and the analog filter cutoff frequency is set to 1000 Hz. Let us define the 4 independent sources of signal as and which stand respectively for the detector dark current, the overall thermal background (due to the sky and the warm optics), the shutter signal (thermal emission of the shutters and downstream warm optics) and the stellar signal. For each star, a few batches of 100 scans ( 1s long) are recorded. In a scan, half of the time is effectively used for stellar interferometric measurements, corresponding to a voltage given by:

where g is the overall gain of the detection process. During the second part of a scan, shutters (see Fig. 1) are closed in the two interferometric arms, and we record the signal:

Ideal background measurements would require fast chopping or nodding on the sky. Since no chopping mode is available yet on IOTA, stellar observations are bracketed by sky observations, pointing towards a direction a few arcseconds away from the star. These sky measurements consist in two batches of 100 scans each, taken roughly 8 mn before and 8 mn after stellar acquisitions, and providing the signals , and:

Fig. 3 shows the various signals recorded: the stellar interferometric raw signal , and for comparison the scans obtained when shutter are closed () and while pointing the sky (). In order to estimate the mean sky level at time of stellar observations t, the shutter signal is used as a common reference. The stellar contribution is estimated through the difference:

The validity of this method obviously depends on the stability of the difference between sky and shutter signals recorded before and after stellar acquisitions. The stability observed is presented in Sect. 3.1.2., showing clearly that observations are detector noise limited, and that we actually sample detector drifts rather than background or sky changes.

Fig. 3a-f. Typical recorded signals (left panels a , c and e ) and corresponding amplitude Fourier spectra (right panels b , d and f in arbitrary units). a  Raw stellar interferogram obtained on . Part of this signal located outside the fringe packet is used to characterize the coupling efficiency of the starlight into the single-mode coupler (Sect. 3.2). Plots c and e show for comparison the temporal scans obtained when shutters are closed, ("shutter signal "), and when tracking a direction a few arcsecs away from the star ("sky signal ").

© European Southern Observatory (ESO) 1999

Online publication: May 6, 1999
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