History of astronomical interferometry

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William Herschel knew as early as 1779 (Herschel 1805) that stars appeared much larger in telescopes than they really were but he did not know why. When Thomas Young demonstrated interference and the wave nature of light unambiguously, this was explained. As he stated in his Bakerian Lecture of 1803: "The proposition on which I mean to insist at present, is simply this, that fringes of colours are produced by the interference of two portions of light", and later: "that homogeneous light, at certain equal distances in the direction of its motion, is possessed of opposite qualities, capable of neutralizing or destroying each other, and of extinguishing the light, where they happen to be united" (Young 1804).

But it was not Young's researches that prompted Herschel to investigate the origin of the spurious diameters of stars. Instead it was the exactly contemporaneous discovery of the first minor planets: Ceres in 1801, Pallas in 1802 and Juno in 1803. Were their apparent diameters as real as those of planets or spurious as for stars? To address this question Herschel conducted an extensive series of experiments in his garden in Slough, examining through his telescope small globules of differing sizes and materials placed in a tree some 800 ft (ca. 244 m) away (Herschel 1805). His observations showed that for the smallest globules the diameters were all spurious and all of the same size. Furthermore, he found that, if just the inner part of the aperture of the telescope were used, the spurious diameters, whether of globules or of stars, were larger. If the whole aperture was employed, the diameters were smaller, and if only an outer annular aperture was used the diameters were smaller still. This experimental discovery that unfilled apertures can be used to obtain high angular resolution remains today the essential basis for interferometric imaging in astronomy (in particular Aperture Masking Interferometry). The theoretical justification of this result came with Airy's analysis of the diffraction pattern of a circular aperture 30 years later (Airy 1835), and it took a further 30 years before the idea of using multiple apertures was developed. In an early study the Reverend W. R. Dawes noted that he had `frequently found great advantage from the use of a perforated whole aperture' and that when observing Venus this produced `a central image of the planet perfectly colourless, and very sharply dened' (Dawes 1866). But it was left to Fizeau, in his submission to the Commission for the Prix Bordin the following year, to remark on `une relation remarquable et n´ecessaire entre la dimension des franges et celle de la source lumineuse' and suggest that by using an interferometric combination of light from two separated slits `il deviendra possible d'obtenir quelques donn´ees nouvelles sur les diametres angulaires de ces astres' (Fizeau 1868).

Steps towards the practical implementation of these techniques for optical astronomy were taken by Michelson, who defined the `visibility' of interference fringes obtained from a source of finite angular size (Michelson 1890) and followed this a year later with the measurement of the angular diameters of Jupiter's satellites (Michelson 1891). Finally, 30 years later, Fizeau's predictions became a reality when the direct interferometric measurement of a stellar diameter was realized by Michelson & Pease (1921) with their 20 ft (ca. 6.1 m) stellar interferometer mounted on the 100 inch Hooker Telescope on Mount Wilson.

1 Interferometric imaging in astronomy
2 Other modern developments
3 References
4 External links

[edit] Interferometric imaging in astronomy

diagram of a radio interferometer

Interferometry provides access to very-high-angular-resolution observations. It also importantly separates the issues of angular resolution and limiting sensitivity. A single mirror of diameter D has an angular resolution 1.22 $λ / D$ , and a collecting area for the flux of photons D², so that there is a well-dened relation between resolution and sensitivity. Since astronomers necessarily study objects beyond their control, it is unlikely that this fixed relationship will be a good match to all but a subset of their observational requirements. It was for this reason that the growth of radio astronomy after 1945 depended so dramatically on the use of interferometric methods. The extreme imbalance between the excellent sensitivity of small filled apertures at long radio wavelengths and their poor angular resolution, which could be many tens of degrees, led naturally to the development of sparse arrays of widely separated telescopes.

In 1946 Ryle and Vonberg (Ryle and Vonberg 1946) constructed a radio analogue of the Michelson interferometer and soon located a number of new cosmic radio sources. The signals from two radio antennas were added electronically to produce interference. Ryle and Vonberg's telescope used the rotation of the Earth to scan the sky in one dimension. Fringe visibilities could be calculated from the variation of intensity with time. Later interferometers included a variable delay between one of the antennas and the detector as shown in the figure at the right.

In the figure radio waves from a source at an angle $θ$ to the vertical must travel a distance $δ l$ further in order to reach the left-hand antenna. These signals are thus delayed relative to the signals received at the right hand antenna by a time $c δ l = c a s i n [θ]$ where $c$ is the speed of the radio waves. The signal from the right hand antenna must be delayed artificially by the same length of time for constructive interference to occur. Interference fringes will be produced by sources with angles in a small range either side of determined by the coherence time of the radio source. Altering the delay time $δ t$ varies the angle at which a source will produce interference fringes. The effective baseline of this interferometer will be given by the projection of the telescope positions onto a plane perpendicular to the source direction. The length of the effective baseline, shown at the bottom of the figure, will be

x = a c o s (θ)

where $a$ is the actual telescope separation.

Technical difficulties delayed the growth of optical interferometry. The human eye is a sensitive detector but is not capable of quantitative photometric assessment of interferometric fringe patterns. When coupled with the need to record data at kHz rates, progress in optical synthesis imaging had to wait for the development of sensitive photon-counting detectors. Furthermore, the limits of mechanical stability had been reached with the 50 ft (ca. 15.2 m) beam interferometer constructed by Michelson and Pease in 1930 (Pease 1931). This method was extended to short-wavelength measurements using separated telescopes by (Johnson, Betz and Towns 1974) in the infrared and by (Labeyrie 1975) in the visible. This demanded micrometre-level metrology of variable optical delay lines, which was not feasible until access to stabilized lasers had become routine. In the 1980s the aperture synthesis technique was extended to visible light and infrared astronomy by the Cavendish Astrophysics Group, providing the first very high resolution images of nearby stars. In 1995 this technique was demonstrated on an array of separate optical telescopes for the first time, allowing a further improvement in resolution, and allowing even higher resolution imaging of stellar surfaces. The same imaging techniques have now been applied at a number of other astronomical telescope arrays, including the Navy Prototype Optical Interferometer, the the CHARA array, the IOTA array and will soon be applied at the VLTI and MRO Interferometer. A detailed description of the development of astronomical optical interferometry can be found here. Impressive results were obtained in the 1990s with COAST and NPOI producing many very high resolution images. Some scientists exaggerated the benefits of combining large diameter (adaptive optics corrected) telescopes for near-infrared interferometry, and this left many astronomers disappointed with new arrays utilizing small numbers of large telescopes which came online in the early 2000s. For details of individual instruments, see the list of astronomical interferometers at visible and infrared wavelengths.

A simple two-element optical interferometer. Light from two small telescopes (shown as lenses) is combined using beam splitters at detectors 1, 2, 3 and 4. The elements creating a 1/4 wave delay in the light allow the phase and amplitude of the interference visibility to be measured, which give information about the shape of the light source.

A single large telescope with an aperture mask over it (labelled Mask), only allowing light through two small holes. The optical paths to detectors 1, 2, 3 and 4 are the same as in the left-hand figure, so this setup will give identical results. By moving the holes in the aperture mask and taking repeated measurements, images can be created using aperture synthesis which would have the same quality as would have been given by the right-hand telescope without the aperture mask. In an analogous way, the same image quality can be achieved by moving the small telescopes around in the left-hand figure - this is the basis of aperture synthesis, using widely separated small telescopes to simulate a giant telescope.

Once these technical considerations had been addressed, all of the principles used at radio wavelengths could be taken over with almost no modification. This included the use of imaging software developed for VLBI at radio wavelengths, which was used to reconstruct the first image from an array of optical telescopes, that of the 50 milliarcsecond binary star Capella (Baldwin et al . 1996). The only signicant differences between the two wavelength regimes is the increased importance of photon shot noise and the small temporal and spatial scales of the atmospheric fluctuations at optical wavelengths. For example, the characteristic time-scale for these fluctuations is measured in milliseconds at optical wavelengths rather than minutes, and the spatial scale is typically smaller than the telescope mirror diameter, whereas at centimetric radio wavelengths this scale can be as large as 20 km. An important consequence of this small spatial scale is that the area of sky over which the atmospheric phase path is constant, the isoplanatic patch, is at most a few arcseconds at visual wavelengths.

[edit] Other modern developments

Between 1950 and 1972, Robert Hanbury Brown and Richard Q. Twiss used optical intensity interferometers to measure the diameters of a large number of stars at visible wavelengths.

Impressive results were obtained in the 1990s, with the Mark III interferometer measuring diameters of 100 of stars and many accurate stellar positions and ISI measuring stars in the mid-infrared for the first time. Additional results include direct measurements of the sizes of and distances to Cepheid variable stars, and young stellar objects.

In the early 2000s single-baseline interferometry became possible with large telescopes, allowing the first measurements of extra-galactic targets. Very primitive imaging has now become technically feasible using large telescopes (using a maximum of 3 VLT telescopes with the AMBER instrument), and it is hoped that by 2008 a useful imaging capability will be available even for extragalactic sources (using e.g. 6 telescopes of the Magdalena Ridge Observatory Interferometer).

Projects are now beginning that will use interferometers to search for extrasolar planets, either by astrometric measurements of the reciprocal motion of the star (as used by the Palomar Testbed Interferometer and the VLTI) or through the use of nulling (as will be used by the Keck Interferometer and Darwin).

[edit] References

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