Astronomical Interferometry
Astronomical interferometry was first imagined by Hippolyte Fizeau in 1868. He proposed to observe stars through two apertures to obtaine interferences that would give an information on the spatial intensity distribution of the source. The resulting resolving power linearly increases with the distance between the apertures and therefore the interference fringes vanish for stars with a diameter proportional to the reciprocal of the baseline of the interferometer. This idea opened the way to a better understanding of most of celestial sources then only know as point-like objects.
This technique is used to combine signals from two or more telescopes to obtain measurements with higher resolution than could be obtained with either telescopes individually. Astronomical interferometry is the basis for astronomical interferometer arrays, which can make measurements of very small astronomical objects if the telescopes are spread out over a wide area. If a large number of telescopes are used a picture can be produced which has resolution similar to a single telescope with the diameter of the combined spread of telescopes. These include radio telescope arrays such as VLA , VLBI , SMA , and more recently astronomical optical interferometer arrays such as COAST , IOTA , resulting in the highest resolution optical images ever achieved in astronomy. In 2001 the "First Fringes" were obtained with the VLT Interferometer (VLTI) at Cerro Paranal in Chile. The VLTI consists in the coherent combination of the four VLT Unit Telescopes and of the four moveable 1.8m Auxiliary Telescopes. Once fully operational, the VLTI will provide both a high sensitivity as well as milli-arcsec angular resolution provided by baselines of up to 200m length. For the VLTI (near infrared and mid-infrared) you can use the follow instrument: VINCI (The VLT INterferometer Commissionning Instrument), AMBER (The near infrared/red VLTI focal instrument), MIDI (The Mid-Infrared interferometric instrument for the VLTI) and PRIMA (The Instrument for Phase-Referenced Imaging and Microarcsecond Astrometry)
Fig.1 The optical train of the VLTI traversed by the two pairs of laser beamstar. In optical, near IR and MID-IR wavelength regimes when we observe a distant object in this way, we see interference fringes. These fringes arise because of the wave nature of light and they contain information about the object being observed. It is important to remember that interferometers record fringes, not images. Than we measure visibility. The angular resolution (in radians) of an interferometer with baseline B is given by
The visibility of fringes is a number between zero and one which measures the fringe contrast (see Figure 3). It is defined as
Fig.2 Comparison of imaging by telescope and interferometer (fringes). 
The visibility curve has a special meaning: it is the Fourier transform of the object's brightness distribution. Thus, to make an image to the sky we need to reserve this process, which means performing an inverse Fourier transform of measured visibilities. To do this accurately for anything but the most simple object requires that you sample the visibility at many different baselines. It is needing good coverage of the (u,v) plane.
Fig.3 Left: examples of fringes with visibilities of 1, 0.5 and 0. Right: visibility as a function of baseline for a resolved star.
Having more telescopes in your interferometer saves time, since it allows you to measure visibilities on several baselines simultaneously (the number of baselines from an array of N telescopes is N(N-1)/2). Allowing the sky to rotate during the night also improves (u,v) coverage and this technique is called earth ritation synthesis.
Fig.4 Binary made of two resolved photometric disks, d=3mas and PA=35 deg. Visibility in (u,v) plane (left) and squared visibility as a function of baseline for different flux ratios 
Fig.5 An example of visibility as a function of baseline for: Uniform disk (green), Binary (black), Gaussian disk (blue) and Multiple object (red).