An optical interferometer samples the wavefronts of light emitted by a source at two or more separate locations and recombines the sampled wavefronts to produce interference fringes. The wavefronts add constructively or destructively, depending on the path difference between the wavefronts, and produce fringes that appear as bright and dark bands, with the bright bands being brighter than the sum of intensities in the two separate wavefronts. A pathlength change in one arm of the interferometer by a fraction of a wavelength causes the fringes to move.

The advantage of interferometry for optical astronomy is that it can provide measurements of stars with a higher angular resolution than is possible with conventional telescopes. Angular resolution is the ability to distinguish accurately two or more points of light that appear close together on the sky.

A stellar long-baseline interferometer consists of an array of several separate telescopes, (at least two) which redirect starlight to a central location where interference fringes are formed.  The angular resolution that the interferometer can achieve depends on the baseline (the telescope separation) while, for a conventional telescope the resolution is restricted by its aperture diameter: the smaller angular separation that can be resolved is proportional to the quantity lambda/d, where lambda is the wavelength of the observation and d is either the  baseline of the interferometer or the diameter of the telescope. To compare sizes the largest existing telescope is 10 m in diameter, whereas the largest interferometer has a baseline of 600 m, giving the latter a factor of 60 increased resolving power.

For ground-based interferometers there are additional aspects to consider in order to to get interferometric fringes: the effects due to the turbulence of the atmosphere and to mechanical vibrations. Both sources of disturbance affect the path difference between the interfering wavefronts, and when this gets too large or varies too quickly during an observation time interferometric fringes cannot be detected. Two conditions need to be met:
 

• The maximum path difference must be smaller than the so called coherence length. This quantity is proportional to lambda squared, so that the smaller the wavelength the stricter the constraints on the path difference that can be tolerated.

• To avoid fringe smearing caused by the wavefronts' paths jittering during an observation time, the variation of the path difference must be kept much smaller than the wavelength of the observation: typically of the order of lambda/10 or lambda/20.

It is a matter of fact that it is less demanding to get interferometric fringes  at long (infrared, far-infrared) wavelengths than at short (visible) wavelengths. In fact, this is the reason why interferometry at radio wavelengths developed much earlier than optical interferometry.

However, the lack of appropriate detector technology to work in the mid-infrared (10 - 20 µm), and the difficulties to operate an instrument which detects its own as well as the sky thermal background (which peaks at these same wavelengths), made the development of near-infrared (2 µm) and visible (0.5 µm) interferometry take the lead over mid-infrared interferometry.

Although several optical interferometers have now been built around the world, the technique is still under development. It has been best used for astrometric purposes, that is for accurate measurement of distances between stars, and for the measurements of stellar diameters. However, by using a sufficient number of telescopes (at least three) and of baselines it is possible to reconstruct images of the observed object with an extremely high angular resolution, and successful application of interferometric imaging techniques to the near infrared wavelengths has already been demonstrated (Baldwin, et al. 1996, Benson et al., 1997 and Hummel, 1998).

Progress in instrumentation has made interferometry a leading technique in the future of high-angular resolution astronomy, and scientific breakthroughs can be expected in fields which span the whole range of astronomical interests.

For more information about basics of stellar interferometry, you can look at the following documents, which are available on the web:

• An Introduction to Interferometry, by Lu Rarogiewicz

• What is an Interferometer?,  from the Keck Interferometer group; sketch and basic description of a stellar interferometer

About MIDI

At the Max-Planck Institut für Astronomie in Heidelberg actions have been taken to exploit our background and technical expertise in mid-infrared astronomy, and the advantages offered by the VLTI for interferometry at these wavelengths, to develop one of the first instruments to produce fringes at mid-infrared wavelengths.
 

The MID-infrared Interferometric instrument (MIDI) is one of the three instruments which will equip the focus of the Very Large Telescope Interferometer (VLTI) (Mariotti 1998). MIDI is a project collaboration between Max-Planck Institut für Astronomie, Heidelberg, Astronomical Institute, Amsterdam, Sterrewacht Leiden, Kapteyn Astronomical Institute Groningen, Observatoire de Paris-Meudon, Observatoire de Nice, Kiepenheuer-Institut Freiburg, Thüringer Landessternwarte Tautenburg.

The actual design of MIDI is optimized for operation at 10 µm and a possible extension to 20 µm is considered. At these wavelength ranges interferometry encounters several disadvantages that however can be coped by means of the unique characteristics of VLTI.
 

• The large apertures of the VLTI Unite Telescopes (8 m diameter) allow to counteract for the high thermal background from both the instrument and the sky. At 10 µm the detection with MIDI will be background limited, and in this case the signal-to-noise ratio scales as S/N ~ D².

 
• The VLTI long baselines make possible to reach good angular resolution also when working at long wavelengths. In fact, since the resolution power goes as lambda/B by using very long baseline (up to 128 m with the UTs) at 10 µm we can reach a resolution of the order of 20 milliarcsecond that in the near-infrared would require baselines 5 times shorter.

 

References

J.E. Baldwin, M.G. Beckett, R.C. Boysen, D. Burns, D.F. Buscher, G.C. Cox, C.A. Haniff, C.D. Mackay, N.S. Nightingale, J. Rogers, P.A.G. Scheuer, T.R. Scott, P.G. Tuthill, P.J. Warner, D.M.A. Wilson, and R.W. Wilson,
The first images from an optical aperture synthesis array: mapping of Capella with COAST at two epochs,
1996, Astron. Astrophys., 306, L13

J.A. Benson, D.J. Hutter, N.M. Elias II, P.F. Bowers, K.J.  Johnston, A.R. Hajian, J.T. Armstrong, D. Mozurkewich, T.A. Pauls, L.J. Rickard, C.A. Hummel, N.M. White, D. Black, and C.S. Denison,
Multichannel optical aperture synthesis imaging of lambda1 URSAE majoris with the Navy prototype optical interferometer,
1997, A. J.,  114, 1221

C.A. Hummel,
Images of spectroscopic binaries with the Navy Prototype Optical Interferometer,
1998,  Proc. SPIE, 3350, 483

J.-M. Mariotti, C. Denise, F. Derie, M. Ferrari, A. Glindemann, B. Koehler, S.A. Leveque, F. Paresce,
M. Schoeller, M. Tarenghi, M. Verola,
VLTI program: a status report,
1998 SPIE, 3350, 800

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