Optical interferometry in astronomy

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Rep. Prog. Phys. 66 (2003) 789–857
PII: S0034-4885(03)90003-0
Optical interferometry in astronomy
John D Monnier
University of Michigan Astronomy Department, 941 Dennison Building, 500 Church Street,
Ann Arbor, MI 48109, USA
E-mail: [email protected]
Received 22 January 2003, in final form 24 March 2003
Published 25 April 2003
Online at stacks.iop.org/RoPP/66/789
Here I review the current state of the field of optical stellar interferometry, concentrating on
ground-based work although a brief report of space interferometry missions is included. We
pause both to reflect on decades of immense progress in the field as well as to prepare for a new
generation of large interferometers just now being commissioned (most notably, the CHARA,
Keck and VLT Interferometers). First, this review summarizes the basic principles behind
stellar interferometry needed by the lay-physicist and general astronomer to understand the
scientific potential as well as technical challenges of interferometry. Next, the basic design
principles of practical interferometers are discussed, using the experience of past and existing
facilities to illustrate important points. Here there is significant discussion of current trends
in the field, including the new facilities under construction and advanced technologies being
debuted. This decade has seen the influence of stellar interferometry extend beyond classical
regimes of stellar diameters and binary orbits to new areas such as mapping the accretion
discs around young stars, novel calibration of the cepheid period-luminosity relation, and
imaging of stellar surfaces. The third section is devoted to the major scientific results from
interferometry, grouped into natural categories reflecting these current developments. Lastly,
I consider the future of interferometry, highlighting the kinds of new science promised by the
interferometers coming on-line in the next few years. I also discuss the longer-term future of
optical interferometry, including the prospects for space interferometry and the possibilities
of large-scale ground-based projects. Critical technological developments are still needed to
make these projects attractive and affordable.
© 2003 IOP Publishing Ltd
Printed in the UK
J D Monnier
1. Introduction
This review will introduce the theory, technique, and scientific goals of optical stellar
interferometry. By combining light collected by widely separated telescopes, interferometrists
can overcome the diffraction-limit of an individual telescope. The angular resolution achieved
by current instruments is indeed astounding, <5 × 10−9 rad (1 milli-arcsecond), as are the
engineering feats of maintaining sub-micron optical stability and coherence over hundreds
of metres of pathlength while controlling polarization and dispersion over broad wavelength
bandpasses. New capabilities are being applied in a wide variety of astrophysics contexts,
including fundamental stellar parameters, novel ways to measure distances to stars, probing
star formation and evolution, direct detection of extrasolar planets, and resolving cores of the
nearest active galactic nuclei (AGN) and brightest quasars. This paper is pointing the way
towards next generation facilities, and I will close with a discussion of efforts to bring the
angular resolution advantages of interferometers into space.
1.1. Scope of review
The history of stellar interferometry spans more than a century, and a proper documentation of
the rich history is beyond the scope of this review (see Lawson (2000a) for a historical overview
of the field). The main purpose of this review will be to summarize the current state of the field of
optical interferometry, including past scientific and engineering lessons, current astronomical
motivations, and future goals and performance expectations. Interested readers may want to
consult earlier reviews which tend to emphasize other topics, especially Shao and Colavita
(1992a) and Quirrenbach (2001). In order to restrict the length, I will concentrate on longbaseline optical interferometry, giving only passing description to diffraction-limited singleaperture experiments (e.g. speckle interferometry, aperture masking, adaptive optics). Further,
I consider ‘optical interferometry’ in a restricted sense to mean the light from the separate
telescopes are brought together using optics, as opposed to heterodyne interferometry whereby
the radiation at each telescope is coherently detected before interference (although I will discuss
briefly the important cases of the intensity interferometer and heterodyne interferometry using
CO2 lasers). In practice then, ‘optical interferometry’ is limited to visible and infrared (IR)
wavelengths, and I will not discuss recent advances in mm-wave and sub-mm interferometry.
1.2. The organization of review
This review is divided into four major sections. The first reviews the basic theory behind optical
interferometry and image reconstruction through a turbulent atmosphere. The second section
explains the basic designs of interferometers and core modern technologies which make them
work, including descriptions of current facilities. Major scientific results are outlined in the
third section. The last section forecasts near-future science potential as well as the long-term
prospects of optical interferometry on the ground and in space.
1.3. Nomenclature
In this review, ‘optical’ does not indicate the visible portion of the electromagnetic spectrum
only, but generally refers to how the light is manipulated (using optics); in the context of
interferometry, this will limit our discussion to wavelengths from the blue (∼0.4 µm) to
the near-IR (1–5 µm) and mid-IR (8–12 µm). Typical angular units used in this paper are
‘milli-arcseconds’, or mas, where an arcsecond is the standard 1/3600 of a degree of angle.
Astronomers often use the magnitude scale to discuss the wavelength-dependent flux density
Optical interferometry in astronomy
(power per unit area per unit bandwidth) from an astronomical source, where the bright star
Vega (α Lyrae) is defined as 0 mag (corresponding to a 10 000 K blackbody); the magnitude
scale is logarithmic such that every factor of 10 brightness decrease corresponds to a flux
‘magnitude’ increase of 2.5 (e.g. a contrast ratio of 10 astronomical magnitudes is a factor
of 104 ). In addition, the unit Jansky (Jy) is often used for measuring the flux density of
astronomical objects, 1 Jy = 10−26 W m−2 Hz−1 .
A number of other basic astronomical units are used herein. The distance between the
Earth and Sun is one astronomical unit, 1 AU 1.5 × 1011 m. Stellar distances are given in
units of parsecs (1 parsec is the distance to a star exhibiting a parallax angle of 1 arcsecond):
written in terms of other common units of length, 1 pc 3.09 × 1016 m ∼ 3.26 light years.
Lastly, I want to alert the reader (in advance) to table 2 which will define all the
interferometer acronyms used throughout the text.
2. Basic principles of stellar interferometry
This section will review the basic principles of stellar interferometry. More detailed discussions
of optical interferometry issues can be found in the recent published proceedings of the
Michelson Summer School, ‘Principles of Long Baseline Stellar Interferometry’ edited by
Lawson (2000b), and the proceedings of the 2002 Les Houches Eurowinter school edited by
Perrin and Malbet (2002); earlier such collections also continue to play an important reference
role (e.g. Perley et al (1986), Lagrange et al (1997)). Other ‘classic’ texts on the subjects of
radio interferometry and optics are Born and Wolf (1965), Goodman (1985) and Thompson
et al (2001).
2.1. Basics of stellar interferometry
The basic principles behind stellar interferometry should be familiar to any physicist, founded
on the wave properties of light as first observed by Thomas Young in 1803. This result is
widely known through Young’s ‘two-slit experiment’, although two-slits were not used in the
original 1803 work.
2.1.1. Young’s two-slit experiment. In the classical set-up, monochromatic light from a distant
(‘point’) source impinges upon two-slits, schematically shown in the left panel of figure 1. The
subsequent illumination pattern is projected onto a screen and a pattern of fringes is observed.
This idealized model is realized in a practical interferometer by receiving light at two separate
telescopes and bringing the light together in a beam-combination facility for interference (this
will be discussed fully in section 3; e.g. see figure 10). The interference is, of course, due to
the wave nature of light á la Huygens; the electric field at each slit (telescope) propagating
to the screen with different relative path lengths, and hence alternately constructively and
destructively interfering at different points along the screen. One can easily write down the
condition for constructive interference; the fringe spatial frequency (fringes per unit angle)
of the intensity distribution on the screen is proportional to the projected slit separation, or
baseline b, in units of the observing wavelength λ (see figure 1). That is,
Fringe spacing ≡ =
Fringe spatial frequency ≡ u =
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Incoming plane waves
Point sources
at infinity
separated by
1/2 the fringe
= Wavelength
Point source
at infinity
2 slits
∆θ = Fringe spacing
λ /b radians
Interference pattern
(Visibility = 1)
2 sine waves
(Visibility = 0)
Figure 1. Young’s two-slit interference experiment (monochromatic light) is presented to illustrate
the basic principles behind stellar interferometry. On the left is the case for a single point-source,
while the case on the right is for a double source with the angular distance being half the fringe
spacing. Note, the interference pattern shown represents the intensity distribution, not the electric
Imagine another point-source of light (of equal brightness, but incoherent with the first)
located at an angle of λ/(2b) from the first source (see right panel of figure 1). The two
illumination patterns are out of phase with one another by 180˚, hence cancelling each other
out and presenting a uniformly illuminated screen. Clearly such an interfering device (an
‘interferometer’) can be useful in studying the brightness distribution of a distant ‘stellar’
object. This application of interferometry was first proposed by Fizeau (1868) and successfully
applied by Michelson to measure the angular diameters of Jupiter’s moons (Michelson 1890,
1891) in 1891 and later (with Pease in 1921) to measure the first angular size of a star beyond
the Sun (Michelson and Pease 1921) (see section 3.1 for further details on the early history of
optical interferometry).
2.1.2. Angular resolution. The ability to discern the two components of a binary star system
is often used to gauge the spatial resolution of an instrument, be it a conventional imaging
telescope or a separated-element interferometer. Classical diffraction theory has established
the ‘Rayleigh Criterion’ for defining the (diffraction-limited) resolution of a filled circular
aperture of diameter D:
This criterion corresponds to the angular separation on the sky when one stellar component
is centred on the first null in the diffraction pattern of the other; the binary is then said to be
resolved. A similar criterion can be defined for an interferometer: an equal brightness binary
Resolution of telescope ≡ telescope = 1.22
Optical interferometry in astronomy
is resolved by an interferometer if the fringe contrast goes to zero at the longest baseline. As
motivated in the last paragraph, this occurs when the angular separation is λ/2b, where b is
the baseline. Hence,
Resolution of interferometer ≡ interferometer =
While these two criteria are somewhat arbitrary, they are useful for estimating the angular
resolution of an optical system and are in widespread use by the astronomical community.
2.1.3. Complex visibility. One can be more quantitative in interpreting the fringe patterns
observed with an interferometer. The fringe contrast is historically called the visibility and,
for the simple (two-slit) interferometer considered here, can be written as
Imax − Imin
Fringe amplitude
V =
Imax + Imin
Average intensity
where Imax and Imin denote the maximum and minimum intensity of the fringes. Hence, the
left and right fringe patterns of figure 1 have visibilities of one and zero, respectively.
The Van Cittert–Zernike theorem (see Thompson et al (2001) for complete discussion and
proof) relates the contrast of an interferometer’s fringes to a unique Fourier component of the
impinging brightness distribution. In fact, the visibility is exactly proportional to the amplitude
of the image Fourier component corresponding to the (spatial) fringe spatial frequency defined
above (u = b/λ rad−1 ). Also, the phase of the fringe pattern is equal to the Fourier phase of
the same spatial frequency component.
The Van Cittert–Zernike theorem can be expressed concisely in mathematical terms.
Consider that the astronomical target emits light at frequency ν over only a very small portion
of the sky with specific intensity Iν (θ, φ), so small that the spherical coordinates θ0 + δθ and
φ0 + δφ can be interpreted as Cartesian coordinates x and y centred around θ0 and φ0 on the
plane of the sky. We can write the interferometer response (amplitude and phase of the fringes)
as the frequency-dependent complex visibility Ṽν (u, v), defined as the Fourier Transform of
= 0) = 1.
the brightness distribution Iν (r ), normalized so that Ṽ (D/λ
dx dy Iν (r )e−2π i((D/λ)·r )
e Vν = δ Vν
λ dx dy Iν (r )
Total specific flux
using the following notation:
r = (x , y ),
projected onto the plane of the sky in units of wavelength λ
≡ the baseline vector D
= (u, v) [Common notation]
Figure 2 shows some simple examples of one-dimensional images and the corresponding
visibility curves. The top panels show the case of an equal binary system, where both
components are unresolved. The periodicity in the visibility-space corresponds to the binary
separation. The middle set of panels is representative of a compact, but resolved, source (such
as a star surrounded by an optically thick dust shell). The small image size means there is
more high spatial frequency information, and this is why the corresponding visibility curve is
non-zero even at high resolution. Lastly, the bottom panels show an image of an unresolved
star (with 10% of the total flux) surrounded by larger-scale structure (this is expected when a
star is surrounded by an optically thin envelope of dust). The large-scale structure (containing
90% of the total flux) can be seen to be ‘resolved’ on short baselines (at low spatial frequency),
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Visibility Curve 1
Image 1
Angle (radians)
Spatial Frequency (radians-1)
Visibility Curve 3
Image 3
Visibility Curve 2
Image 2
Spatial Frequency (radians-1)
Spatial Frequency (radians-1)
Figure 2. This figure shows simple one-dimensional images and their corresponding visibility
curves. The left panels are the images while the right panels correspond to the Fourier amplitudes,
i.e. the visibility amplitudes. Note that ‘large’ structure in image-space result in ‘small’ structure
in visibility-space.
while the point-source remains unresolved out to the highest spatial frequency. Note that the
visibility plateaus at 0.10, corresponding to the fraction of the total flux which is left unresolved.
This is easy to understand since the Fourier Transform is linear; that is, the (complex) visibility
of a point-source and extended structure is equal to the visibility of the point-source plus the
visibility of the extended structure separately. This property of linearity is very helpful in
interpreting simple visibility curves.
Most astronomical objects are not one-dimensional, and the two-dimensional space of
spatial frequencies is called the Fourier Plane, or the (u, v) plane, named after the (u, v)
coordinates defined in equation (7). Further, in general we must consider both the visibility
amplitude and the visibility phase. For example, consider the equal binary system depicted in
figure 3. The complex visibility can be easily written by choosing the origin midway between
the two components. Note the abrupt phase jump when the visibility amplitude goes through
a null. These discontinuities are smoothed out when the two components are not precisely
2.2. Atmospheric problems
An incoming plane wave from a stellar source is corrupted as it propagates through the turbulent
atmosphere. Variations in the column density of air along different paths cause the effective
pathlength to vary, introducing wavefront distortion. If these distortions become a significant
Optical interferometry in astronomy
Figure 3. This figure shows the complex visibility for an equal binary system in the two-dimensional
(u, v) plane. With the above choice for the phase centre, the Fourier phases can be represented
simply. Notice the abrupt phase jumps when the visibility amplitude goes through a null. This figure
is reproduced through the courtesy of the NASA/Jet Propulsion Laboratory, California Institute of
Technology, Pasadena, California (Monnier 2000).
fraction of a wavelength across the aperture of a telescope, the size of the image formed will not
be diffraction-limited by the primary mirror, but rather by the coherence scale of the incoming
wavefront. The transverse distance over which one expects rms pathlength difference to be
λ/2.4 has been defined as the Fried parameter and is denoted by r0 (λ) (Fried 1965); hence
telescope apertures larger than r0 (λ) can expect significant degradation of image quality (when
observing at wavelength λ) due to atmospheric effects. In fact, for an r0 diameter circular
patch, the rms phase error is ∼1.03 rad. At λ = 500 nm, r0 is typically 10 cm (toward zenith)
at average observing sites and hence even small telescopes cannot be used at their diffractionlimit in the visible. In such cases, the observed angular size of a point-source will be determined
entirely by r0 (λ) at a given wavelength, and is known as the seeing disc size, seeing (λ). The
Kolmogorov theory of turbulence (Kolmogorov 1961) predicts that r0 (λ) ∝ λ6/5 , and hence
the seeing size, seeing (λ) ∝ λ/r0 (λ) ∝ λ−1/5 , is only weakly dependent on the wavelength
(Fried 1965). An example of the phase delays associated with a snapshot of Kolmogorov
turbulence can be seen in figure 4 for 12 m square, corresponding roughly to the size of the
largest telescopes today (e.g. the Keck telescopes).
Another consequence of turbulence is that the image distortion varies across the sky,
although stars located close together suffer similar seeing effects. The angle over which image
distortions are correlated is called the ‘isoplanatic’ angle, and is only a few arcseconds in
the visible and about an arcminute in the near-IR. This angle is determined by the vertical
distribution of the turbulence—obviously low-level turbulence would induce correlated image
distortions over larger sky angles than the same turbulent layer located higher up in the
atmosphere. The isoplanatic angle is a critical parameter for the field-of-view of adaptive
optics systems which actively sense and correct for atmospheric turbulence in real-time (RT)
using a deformable mirror.
Another important atmospheric diagnostic is the coherence time, t0 . Typically, one
assumes a ‘frozen’ turbulence model in which the atmospheric density perturbations are
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Figure 4. This figure shows a typical realization of Kolmogorov turbulence (r0 = 50 cm at
λ = 2.2 µm); each solid contour line represents λ/2 of wavefront distortion. Some areas
of the aperture show coherent areas larger than r0 , and some much smaller; r0 is a statistical
property of atmospheric turbulence and wavefront perturbations occur over a wide range of
assumed constant over the time it takes wind to blow them across a given aperture (also known
as Taylor’s hypothesis of frozen turbulence). This motivates a convenient estimate for the
coherence time: t0 (λ) ≡ r0 (λ)/vwind , where vwind is the wind speed. At most sites, wind speeds
are ∼10 m s−1 and so t0 ∼ 10 ms at 500 nm. Further discussion of atmospheric turbulence and
degradation of astronomical images can be found in Kolmogorov (1961), Roddier (1981), and
Roddier et al (1982).
These parameters, r0 (λ) and t0 (λ), are extremely important for the design of an
interferometer, because the value of r0 limits the useful size of the collecting aperture, while t0
limits the coherent integration time. Both of these are crucial for predicting an interferometer
sensitivity to faint objects and much debate surrounds the best estimates for these parameters
at various sites (e.g. Dyck and Howell (1983), Roddier et al (1990), ten Brummelaar (1994),
Treuhaft et al (1995)). This topic is revisited in section 3.4 when I discuss the limiting
magnitude of current interferometers.
It is well-known that r0 (λ) and t0 (λ) depend greatly on the observing site, and we now
consider the unique seeing conditions of Mauna Kea, Hawaii, as an example. The coherence
scale is unusually long due to the highly laminar flow of the Pacific winds over the peak of the
mountain (elevation 4200 m); r0 usually lies between 10 and 40 cm at 500 nm (Wizinowich
1999). However, the fast winds of the overhead jet stream result in very short coherence
times: t0 between 1.5 and 10 ms (Wizinowich 1999). It should be emphasized that seeing is
notoriously difficult to characterize due to large variations in time (both on short timescales as
well as seasonal ones) as evidenced by the large range of r0 (λ = 500 nm) and t0 (λ = 500 nm)
values just given.
2.2.1. Atmospheric phase errors. The fluctuating amount of integrated atmospheric
pathlength above each telescope introduce wavefront time delays which show up as phase shifts
in the measured fringes in an interferometer, as illustrated in figure 5. In this figure, an optical
interferometer is represented again by a Young’s two-slit experiment, as discussed earlier in this
Optical interferometry in astronomy
Point source
at infinity
Incoming plane waves
Figure 5. Atmospheric time delays or phase errors at telescopes cause fringe shifts, as can be seen
through analogy with Young’s two-slit experiment. This figure is reproduced through the courtesy
of the NASA/Jet Propulsion Laboratory, California Institute of Technology, Pasadena, California
(Monnier 2000).
section. The spatial frequency of these fringes is determined by the distance between the slits
(in units of the wavelength of the illuminating radiation). However if the pathlength above oneslit is changed (e.g. due to a pocket of warm air moving across the aperture), the interference
pattern will be shifted by an amount depending on the difference in pathlength of the two
legs in this simple interferometer. If the extra pathlength is half the wavelength, the fringe
pattern will shift by half a fringe, or π rad. The phase shift is completely independent of the
slit (telescope) separation, and only depends on slit-specific (telescope-specific) phase delays.
The most obvious impact of atmospheric phase delays is that the assumptions of the van
Cittert–Zernike theorem no longer apply, and that the measure fringe phase can no longer be
associated with the Fourier phase of the sky brightness distribution (the fringe amplitude retains
its original meaning, since phase changes do not change the measured fringe amplitudes for
short exposures). The corruption of this phase information has serious consequences, since
imaging of non-centrosymmetric objects rely on the Fourier phase information encoded in this
intrinsic phase of interferometer fringes. Without this information, imaging cannot be done
except for simple objects such as discs or round stars. Fortunately, a number of strategies have
evolved to circumvent these difficulties.
2.2.2. Phase referencing. Possible methods for recovering this phase information using phase
referencing techniques are discussed in chapter 9 (written by A Quirrenbach) of ‘Principles
of Long Baseline Stellar Interferometry’ (Lawson 2000b). Few scientific results have resulted
from phase referencing techniques to date, but this is expected to change over the coming
decade as sophisticated new instruments are being deployed. Here, I mention a few of the
most promising methods:
(i) Nearby sources. If a bright point reference source (or source with well-known structure)
lies within an isoplanatic patch (see Quirrenbach (2000)), then its fringes will act as a
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probe of the atmospheric conditions. By measuring the instantaneous phases of fringes
from the bright reference source, one can correct the corrupted phases on a nearby ‘target’
source. This has been applied to narrow-angle astrometry where fringe phase information
is used for determining precise relative positions of nearby stars (Shao and Colavita 1992b,
Colavita et al 1999, Lane et al 2000a); see figure 31 for some preliminary results published
by the Palomar Testbed Interferometer. While it would be very valuable to use an artificial
guide star for phase-referencing a long-baseline interferometer, current laser beacons are
too spatially extended (resolved) to produce interferometric fringes.
(ii) Monitoring. In the millimetre and sub-millimetre, phase shifts caused by fluctuations
in atmospheric water vapour column density can be monitored by observing its line
emission. This information can be used to phase-compensate the interferometer, allowing
longer coherent integrations and accurate fringe phase determination on the target
(see Wiedner (1998), and references therein). In the mid-IR, strategies to actively
monitor ground-level turbulence using temperature sensors are being explored by the
Infrared Spatial Interferometer group (Short et al 2003) at Mt Wilson motivated by
recent atmospheric studies (e.g. Bester et al (1992)). Townes (2002) recently proposed
that RT-monitoring of Rayleigh or Raman backscattering might be used to correct for
atmospheric column density variations in the context of optical interferometers, but this
method has not yet been validated.
(iii) Multi-wavelength. Another possibility is to observe a target at multiple wavelengths and
to use data from one part of the spectrum to calibrate another. For example, one might
use fringes formed by the continuum emission to phase reference a spectral line (e.g.
Vakili et al (1997)). To use this method, one must assume knowledge about the brightness
distribution at one of the wavelengths being used.
Currently, phase referencing is not possible with most current beam combiners in
operation, either due to low spectral resolution or limited field-of-view. In order to recover
phase information, one must make use of the closure phases.
2.2.3. Closure phases. Consider figure 6 in which a phase delay is introduced above
telescope 2 of a 3-telescope array. As discussed in the last section, this additional delay
causes a phase shift in the fringe detected between telescopes 1 and 2. Note that a phase shift
is also induced for fringes between telescopes 2 and 3; however, this phase shift is equal and
opposite to the one for telescopes 1 and 2. Hence, the sum of three fringe phases, between
1–2, 2–3, and 3–1, is insensitive to the phase delay above telescope 2. This argument holds for
arbitrary phase delays above any of the three telescopes. In general, the sum of three phases
around a closed triangle of baselines, the closure phase, is a good interferometric observable;
that is, it is independent of telescope-specific phase shifts induced by the atmosphere or
The idea of closure phase was first introduced to compensate for poor phase stability in
early radio very long-baseline interferometry (VLBI) work (Jennison 1958). Application at
higher frequencies was first mentioned by Rogstad (1968), but only much later carried out in
the visible/IR through aperture masking experiments (Baldwin et al 1986, Haniff et al 1987,
Readhead et al 1988, Haniff et al 1989). Currently three separate-element interferometers have
succeeded in obtaining closure phase measurements in the visible/IR, first at COAST (Baldwin
et al 1996), soon after at NPOI (Benson et al 1997), and most recently at IOTA (Traub 2003).
How can these closure phases be used to figure out the Fourier phases which are needed to
allow an image to be reconstructed? Each closure triangle phase can be thought of as a single
linear equation relating three different Fourier phases (assuming none of the baselines are
Optical interferometry in astronomy
Φ(1-2) = ΦΟ(1-2) + [φ(2)-φ(1)]
Φ(2-3) = ΦΟ(2-3) + [φ(3)-φ(2)]
Φ(3-1) = ΦΟ(3-1) + [φ(1)-φ(3)]
Closure = Φ (1-2)+Φ (2-3)
(1-2-3) +ΦΟ(3-1)
Figure 6. This figure explains the principle behind closure phase analysis. Phase errors introduced
at any telescope causes equal but opposite phase shifts, cancelling out in the closure phase (figure
after Readhead et al (1988)). This figure is reproduced through the courtesy of the NASA/Jet
Propulsion Laboratory, California Institute of Technology, Pasadena, California (Monnier 2000).
identical), which we desire to solve for; hence, we must count the number of linear equations
N and compare to the number of unknowns. For N telescopes, there are ‘N choose
N 3 = ((N )(N − 1)(N − 2))/(3)(2), possible closing triangles. However, there are only
= ((N )(N − 1))/2 independent Fourier phases; clearly not
all the closure phases can be
independent. The number of independent closure phases is only N 2−1 = ((N − 1)(N − 2))/2,
equivalent to holding one telescope fixed and forming all possible triangles with that telescope
(as discussed by Readhead et al (1988)). The number of independent closure phases is always
less than the number of phases one would like to determine, but the per cent of phase information
retained by the closure phases improves as the number of telescopes in the array increases.
Table 1 lists the number of Fourier phases, closing triangles, independent closure phases, and
recovered percentage of phase information for telescope arrays of 3–50 elements. For example,
approximately 90% of the phase information is recovered with a 21 telescope interferometric
array (e.g. Readhead et al (1988)). As discussed in the next section, this phase information
can be coupled with other image constraints (e.g. finite size and positivity) to reconstruct the
source brightness distribution.
In addition to the mathematical (linear algebra) interpretation of closure phases, there are
a few other important properties worth noting.
• For sources with point-symmetry (otherwise known as centro-symmetry), all the closure
phases are either 0˚ or 180˚. It is easy to prove this by imagining the image-centre (‘phasecentre’) at the point of centro-symmetry.
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Table 1. Phase information contained in the closure phases alone.
Number of
Number of
Fourier phases
Number of
closing triangles
Number of independent
closure phases
Percentage (%) of
phase information
1 330
2 925
19 600
• Closure phases are not sensitive to an overall translation of image. A translation is
indistinguishable from atmospheric phase delays for any given closing triangle.
• The closure phases are independent of telescope-specific phase errors, however non-zero
closure phases from a point-source can result from having non-closing triangles and phase
delays in the beam combiner (e.g. for a 3-telescope pair-wise beam combiner).
2.3. Image reconstruction
While very few images have been made by today’s optical interferometers, new telescope arrays
are now being commissioned which will make true imaging interferometry straightforward.
Because these new imaging capabilities are likely to have significant impacts over the next
decade, I wish to review the basic principles of apertures synthesis imaging. However, I will
restrain myself from excessive elaboration here, and instead refer the interested reader to the
extensive radio interferometry literature, especially regarding ‘VLBI’.
While modelling visibility and closure phase data with simple models is useful, one
would like to make an image unbiased by theoretical expectations. Since any image can
be alternatively represented by its Fourier components, the collection of all ‘interesting’
components can allow the interferometric data to be inverted, thus reconstructing an estimation
of the image brightness distribution. The collection of a large number of Fourier components
is greatly aided by increasing the number of telescopes, since independent combinations of
telescopes increase with the number of telescopes to the second power ((N )(N − 1))/2; see
last section.
With a large number of measurements, images of arbitrary complexity should be attainable
using visible/IR interferometers and reliable closure phase measurements. The importance of
‘filling up’ the (u, v) plane with measurements when imaging is discussed more fully in
section 2.4.3. The next subsections will discuss strategies currently employed based on the
techniques of VLBI in the radio.
2.3.1. Guiding principles. The goals of an image reconstruction procedure can be stated
quite simply: find an image which fits both the visibility amplitudes and closure phases within
experimental uncertainties. However in practice, there are an infinite number of candidate
images which satisfy these criteria, because interferometric data is always incomplete and
noisy. Furthermore, the closure phases cannot be used to unambiguously arrive at Fourier
phase estimates as stated above, even under ideal noise-free conditions.
Additional constraints must be imposed to ‘select’ an image as the best-estimate of the
true brightness distribution (to ‘regularize’ this ill-posed inverse problem). These constraints
introduce correlations in the Fourier amplitudes and phases, and essentially remove degrees of
freedom from our inversion problem. Some of the most common (and reasonable) constraints
Optical interferometry in astronomy
are described below.
• Limited field-of-view. This constraint is always imposed in aperture synthesis imaging,
even for a fully-phased array (e.g. VLA). Limiting the field-of-view introduces
correlations in the complex visibility in the (u, v) plane. This is a consequence of the
convolution theorem, a multiplication in image-space is equivalent to a convolution in the
corresponding Fourier-space.
• Positive-definite. Since brightness distributions cannot be negative, this is a sensible
constraint (although not appropriate in some cases, such as for reconstruction of
Stokes/polarization components or imaging spectral line absorption). While clearly
limiting the range of ‘allowed’ complex visibilities, there are few obvious, intuitive
effects in the Fourier-plane; one is that the visibility amplitude is maximum at zero
spatial frequency. The maximum entropy method (MEM) (see section 2.3.3) naturally
incorporates this constraint.
• ‘Smoothness’. MEM, for instance, selects the ‘smoothest’ image consistent with the data
(see section 2.3.3 for more discussion of MEM).
• A priori information. One can incorporate previously known information to constrain the
possible image reconstructions. For instance, a low resolution image may be available
from a single-dish telescope. Another commonly encountered example is point-source
embedded in nebulosity; one might want the reconstruction algorithm to take into account
that the source at the centre is point-like from theoretical arguments.
For a phased interferometric array (e.g. the Very Large Array (VLA)) where the Fourier
phases are directly measured (avoiding the need for closure phases), one can use a number
of aperture synthesis techniques to produce an estimate of an image based on sparsely
sampled Fourier components. These procedures basically remove artefacts, i.e. sidelobes,
of the interferometer’s point-source response arising from incomplete sampling of the (u, v)
plane. These procedures do not incorporate closure phases, but work by inverting the Fourier
amplitudes and phases to make an image. A brief explanation of the popular algorithms
CLEAN and MEM follow with additional references for the interested reader (see Perley et al
(1986) for essays on these topics aimed at radio astronomers).
2.3.2. CLEAN. Originally described by Högbom (1974), CLEAN has been traditionally the
most popular algorithm for image reconstruction in the radio because it is both computationally
efficient and intuitively understandable. Given a set of visibility amplitudes and phases over a
finite region of the Fourier plane, the ‘true’ image can be estimated by simply setting all other
spatial frequencies to zero and taking the (inverse) Fourier transform. As one might expect, this
process leads to a whole host of image artefacts, most damaging being positive and negative
‘sidelobes’ resulting from non-complete coverage of the Fourier plane; we call this the ‘dirty
map’. The unevenly filled Fourier plane can be thought of as a product of a completely sampled
Fourier plane (which we desire to determine) and a spatial frequency mask which is equal to
1 where we have data and 0 elsewhere. Since multiplication in Fourier space is identical to
convolution in image space, we can take the Fourier transform of the spatial frequency mask
to find this convolving function; we call this the ‘dirty beam’. Now the image reconstruction
problem can be recast as a ‘deconvolution’ of the dirty map with the dirty beam.
The dirty map is CLEANed by subtracting the dirty beam (scaled to some fraction of
the map peak) from the brightest spot in the dirty map. This removes sidelobe structure and
artefacts from the dirty map. Repeating this process with dirty beams of ever decreasing
amplitudes leads to a series of delta-functions which, when combined, fit the interferometric
data. For visualization, this map of point-sources is convolved with a Gaussian function whose
J D Monnier
FWHM values are the same as the dirty beam; this removes high spatial resolution information
beyond the classic ‘Rayleigh’ criterion cutoff. One major weakness with CLEAN is that this
smoothing changes the visibility amplitudes, hence the CLEANed image no longer strictly
fits the interferometric data, especially the spatial frequency information near the diffractionlimit. Another weakness is that CLEAN does not directly use the known uncertainties in the
visibility data, and hence there is no natural method to weight the high signal-to-noise ratio
(SNR) data more than the low SNR data during image reconstruction. Further discussion
of various implementations of CLEAN can be found in Clark (1980), Cornwell (1983) and
Schwab (1984), and chapter 7 of Perley et al (1986) by T Cornwell.
2.3.3. MEM. The MEM makes better use of the highest spatial frequency information
by finding the smoothest image consistent with the interferometric data. While enforcing
positivity and conserving the total flux in the frame, ‘smoothness’ is estimated here by a
global scalar
quantity S, the ‘entropy’. If fi is the fraction of the total flux in pixel i, then
S = − i fi ln(fi /Ii ) after the thermodynamic quantity; Ii is known as the image prior and
must be specified by the user. The MEM map fi will tend toward Ii when there is little (or
noisy) data to constrain the fit. Often Ii is assumed to be a uniformly bright background,
however one can use other image priors if additional information is available, such as the
overall size of the source which may be known from previous observations.
Mathematically, MEM solves the multi-dimensional (N = number of pixels) constrained
minimization problem which only recently has become computationally realizable on desktop
computers. Maintaining an adequate fit to the data (χ 2 ∼ number of degrees of freedom),
MEM reconstructs an image with maximum S. MEM image reconstructions always contain
some spatial frequency information beyond the diffraction-limit in order to keep the image as
‘smooth’ as possible consistent with the data. Because of this, images typically have maximum
spatial resolution a few times smaller than the typical Rayleigh-type resolution encountered
with CLEAN (‘super-resolution’). Further discussions of MEM and related Bayesian methods
can be found in Gull and Skilling (1983), Skilling and Bryan (1984), Narayan and Nityananda
(1986), Sivia (1987) and Pina and Puetter (1992).
Unfortunately, MEM images also suffer from some characteristic artefacts and biases.
Photometry of MEM-deconvolved images is necessarily biased because of the positivity
constraint; any noise or uncertainty in the imaging appears in the background of the
reconstruction instead of the source, systematically lowering the estimated fluxes of compact
sources. Also, fields containing a point-source embedded in extended emission often show
structure reminiscent of Airy rings, the location of the rings being influenced by the wavelength
of the observation and not inherent to the astrophysical source. Fortunately, these imaging
artefacts are greatly alleviated for asymmetric structures, when closure phases and not the
visibility amplitudes play a dominant role in shaping the reconstructed morphology.
2.3.4. Including closure phase information. The above algorithms were originally designed
to use Fourier amplitudes and phases, not closure phases. In order to use these algorithms,
one has to come up a way to estimate the Fourier phases, when only the closure phases
are available. Early image reconstruction algorithms incorporated closure phase information
by using an iterative scheme (Readhead and Wilkinson 1978, Thompson et al 1986) . The
following steps summarize this process:
(i) Start with a Fourier ‘phase model’ based on either prior information or setting all phases
to zero.
Optical interferometry in astronomy
(ii) Determine candidate phases by using some values from the ‘phase model’ and enforcing
all the (self-consistent) closure phase relations (see section 2.2.3).
(iii) Using CLEAN or MEM, perform aperture synthesis mapping on the given visibilities and
candidate phases. At this stage, image constraints such as positivity and/or finite support
are applied.
(iv) Use this image as a basis for a new ‘phase model’.
(v) Go to step (ii) and repeat until the process converges to a stable image solution.
Cornwell and Wilkinson (1981) introduced a modification of the above scheme by
explicitly solving for the telescope-specific errors as part of the reconstruction step. Hence
the measured (corrupted) Fourier phases are fit using a combination of intrinsic phases (which
are used for imaging using CLEAN/MEM) plus telescope phase errors. In this scheme, the
closure phases are not explicitly fit, but rather are conserved in the procedure since varying
telescope-specific errors cannot change any of the closure phases. Figure 7 shows a flow
diagram for this procedure; thoughtful consideration is required in order to fully understand
the power and elegance of self-calibration, affectionately known as ‘self-cal’.
Self-calibration works remarkably well for large number of telescopes, but requires
reasonably high SNR (SNR 5) complex visibilities. Once the SNR decreases below this
point, the method completely fails. The conceptualization of solving for telescope-specific
errors, while useful for the radio, is not applicable for visible/IR interferometry where the
good observables are the closure phases themselves, not corrupted Fourier phases. This is
because the timescale for phase variations in the visible/IR is much less than a second, as
opposed to minutes/hours in the radio.
Of course, the self-calibration iteration loop can be sidestepped altogether by fitting
directly to all the data, the visibility amplitudes and closure phases, using MEM or some other
models intrinsic Fourier phases
plus telescope errors
ij = Φ
ij – (φi – φj)
telescope errors
Generate Fourier phases consistent with
closure phases & begin with initial trial image
Calculate complex visibility (amplitudes
and phases) of trial image
Adjust telescope errors so measured Fourier
phases are best fit by combination of trial image
phases plus telescope errors
Correct trial phases based on new estimates of
telescope errors, and map using CLEAN/MEM
If not converged, use this new map
as the next trial image
Figure 7. This is a flow diagram for an incorporating closure phase information into CLEAN/MEM
aperture synthesis imaging algorithms based on the ‘self-calibration’ procedure of Cornwell and
Wilkinson (1981). This figure is reproduced through the courtesy of the NASA/Jet Propulsion
Laboratory, California Institute of Technology, Pasadena, California (Monnier 2000).
J D Monnier
regularization scheme. This would have the added advantage of allowing all the measurement
errors to be properly addressed, theoretically resulting in the optimal image reconstruction.
Buscher (1994) suggested this approach, but there has been little demonstrated progress in this
method to date. I anticipate revived activity as more interferometers with ‘imaging’ capability
begin to produce data.
2.3.5. Speckle interferometry. Another related interferometric technique which permits
diffraction-limited observation through a turbulent atmosphere using a single filled-aperture
telescope is ‘speckle interferometry’, the promise of which was first realized by Labeyrie
(1970) in 1970. In section 2.2, I claimed that observed angular size of a point-source will be
determined entirely by r0 at a given wavelength, and is known as the seeing disc size, seeing .
However, this is only true for a long-exposure image. A single short-exposure image of a star
actually consists of a network of small ‘speckles’ extending over seeing .
In the original formulation of speckle interferometry, short exposures of an astrophysical
object are made to freeze this ‘speckling’ induced by the turbulent atmosphere. The amount
of high-resolution structure in the speckle pattern, as quantified by its power spectrum, is a
measure of two things: (1) the quality of the atmospheric seeing, and (2) the high resolution
structure in the object of interest. Observing a nearby point-source star allows the calibration of
the seeing contribution and thus the extraction of interferometric visibility measurements out
to the diffraction-limit of the telescope (i.e. the longest baseline). In analysing this situation,
one can think of many virtual sub-apertures (with size equal to the coherence length r0 ) spread
across the full telescope, with fringes forming between all the sub-aperture pairs. After the
original formulation by Labeyrie, it was discovered that the Fourier phases could also be
estimated from such data (e.g. Knox and Thompson (1974), Weigelt (1977)).
Speckle interferometry data is often reduced using the ‘bispectrum’, which permits a direct
inversion from the estimated Fourier amplitudes and phases. The bispectrum B̃ij k = Ṽij Ṽj k Ṽki
is formed through triple products of the complex visibilities around a closed triangle, where
ijk specifies the three aperture locations on the pupil of the telescope. One can see the
bispectrum is a complex quantity, and that the bispectrum phase is identical to the closure
phase. Interestingly, the use of the bispectrum for reconstructing diffraction-limited images
was developed independently (Weigelt 1977, Hofmann and Weigelt 1993) of the closure phase
techniques, and the connection between the approaches realized only later (Roddier 1986,
Cornwell 1987).
2.4. Other important considerations
2.4.1. Coherency. The tolerance for matching pathlengths in an interferometer depend on
the desired spectral bandwidth. In the limit of monochromatic light, such as for a laser,
interference will occur even when pathlengths of an interferometer are highly mismatched.
For broadband (‘white’) light, the number of fringes in an interferogram is equal to the inverse
of the fractional bandwidth: Nfringes ∼ λ/λ. Hence, for broad band observations (∼20%
bandwidth) interference is only efficient if the pathlengths are matched to within a wavelength
or so—a stringent requirement.
2.4.2. Field-of-view. One consequence of the short coherency envelope for broadband
observations is a limitation on the field-of-view. Bandwidth-smearing, as it is called, limits
the field-of-view to be equal to the fringe-spacing × the number of fringes in the coherency
envelope (see last subsection), FOV ∼ λ/Baseline × (λ/λ) rad. This ‘field-of-view’ is thus
baseline-dependent, leading to confusing interpretations of data for extended sources. While
Optical interferometry in astronomy
this effect can be modelled, it should be avoided by using a spectrometer to limit the bandwidth
of individual observing channels.
Another common limitation of the field-of-view is the primary beam of an individual
telescope. For most kinds of beam-combiners, flux outside the diffraction-limited beam is
rejected (spatial filtering is described in detail in section 3.5.1). For most astronomical objects
observed by interferometers, this is not a serious problem since long integrations by (lowresolution) individual telescopes can be used to confirm that no significant flux arises from
outside the primary beam. For an imaging interferometer, one would like to use narrow
enough bandwidths so that bandwidth-smearing (on the longest baselines) is small enough
so that the entire primary beam can be mapped. The requirement for this is approximately:
λ/λ ∼ longest baseline/telescope diameter.
Of course, having a wide field-of-view would be useful for many studies, such as measuring
proper motions of stars at the galactic centre. As discussed later, a wide field-of-view (beyond
the primary beam) can only be achieved in a so-called Fizeau combiner. The Large Binocular
Telescope Interferometer is the only interferometer currently being built which will have this
unique and potentially very powerful faculty.
2.4.3. Filling the (u, v) plane. The ability to make an image depends most strongly on the
filled fraction of the (u, v) plane. Recall that the visibility amplitude and phase measured by
an interferometer is directly related to a single component of the Fourier Transform of the
object brightness distribution. If the object brightness is specified on coordinates of Right
Ascension (pointing East) and Declination (pointing North), then the reciprocal Fourier space
has axes referred to as (u, v). Following astronomical notation, the positive u-axis typically
points to the left on a diagram, just as right-ascension coordinates increase towards the left
For a fixed geometry of telescope locations, the Fourier coverage varies as the star rises
and sets, and as a function of the star’s declination and the interferometer’s latitude. Figure 8
shows the Fourier coverage of three actual interferometers (number of telescopes 3, 6, and
21 for IOTA, CHARA, and Keck aperture masking, respectively) for a declination 45˚ object
spanning 3 h before and after transit (assuming monochromatic light). In aperture masking,
the pupil plane of a single telescope is split up into sub-pupils which are allowed to combine,
just like a long-baseline interferometer. It is obvious that the coverage increases rapidly with
number of telescopes. It is not so obvious that the number of closure phases/triangles also
rapidly increases with array size (equivalent to filling up the hyper-volume (u1 , v1 , u2 , v2 )
with closure triangles; see table 1). Tuthill and Monnier (2000) studied how imaging fidelity
and dynamic range are affected by differing amounts of Fourier coverage using real data.
Obviously for imaging it is absolutely critical to collect as much coverage as possible, and
suitable array design is further discussed in section 3.2.
3. Basic designs of stellar interferometers
3.1. Brief historical overview
Here I give only a brief historical overview of progress in stellar interferometry drawn partially
from the review by Lawson (2000a); I refer the interested reader to the above paper for more
Modern interferometry can be traced back to 19th century France. Hippolyte Fizeau first
outlined in 1868 the basic concept of stellar interferometry, how interference of light could be
used to measure the sizes of stars. The first attempts to apply this technique, akin to modern-day
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An IOTAConfiguration)
Keck Aperture Mask
Position (metres, North)
Position (metres, North)
Position (metres, North)
10 15
Position (metres East)
IOTA Fourier Coverage
Position (metres East)
CHARA Fourier Coverage
Keck Fourier Coverage
Position (metres East)
North Baseline (m)
North Baseline (m)
North Baseline (m)
East Baseline (m)
-300 -200 -100 0 100 200 300
East Baseline (m)
East Baseline (m)
Figure 8. Example of (u, v) plane coverage for different interferometers. The top panels show the
interferometer array configurations, while the bottom panels show the corresponding (u, v) plane
coverage. For the IOTA and CHARA interferometers, I have assumed a source at 45˚ declination
observed for 3 h both before and after transit. The right-most panels show instantaneous ‘snapshot’
coverage for an optimized 21-telescope array, a geometry actually used in the Keck aperture masking
experiment (Tuthill et al 2000c). Note that the circles in the top plot are not to the same scale as
the individual telescope diameters but have been enlarged.
‘aperture masking’, were carried out by E Stéphan soon thereafter, although the telescopes of
that time had insufficient resolution to resolve even the largest stars.
Albert Michelson developed a more complete mathematical framework for stellar
interferometry in 1890; while apparently Michelson was unaware of Fizeau’s earlier work,
more historical investigation is needed to establish this definitively. Along with Pease,
Michelson (Michelson and Pease 1921) eventually succeeded in measuring the diameter of
α Orionis (Betelgeuse) in 1920–1921 using the Mt Wilson 100 in. telescope (following earlier
measurements of Jupiter’s moons; Michelson (1890, 1891)). Interestingly, Michelson needed
a baseline longer than 100 in. in order to resolve Betelgeuse (uniform disc (UD) diameter
∼47 mas), and acquired one by installing a 20-foot interferometer beam on the Cassegrain cage
as illustrated in figure 9, reproduced here from their original paper. Following the success of the
20-foot interferometer, Pease (with Hale) constructed a 50-foot interferometer (on Mt Wilson,
but separate from the 100 in. telescope); although some results were reported, this experiment
was not very successful. Due to its generally outstanding atmospheric conditions, Mt Wilson
continued to be a choice site for interferometry projects, subsequently hosting the Mark III,
ISI, and CHARA interferometers.
Following the disappointing results from the 50-foot interferometer, it would be decades
before significant developments inspired new activity in the optical arena. Meanwhile,
advances in radar during World War II spurred rapid development of radio interferometry.
We refer the reader to Thompson et al (2001) for a discussion of the development of radio
interferometry beginning with the first radio interferometer built by Ryle and Vonberg in 1946.
Optical interferometry in astronomy
Figure 9. This diagram from Michelson and Pease (1921, figure 1) illustrates how a 20-foot
interferometer beam was installed on the Mt Wilson 100 in. telescope in order to create, for the
first time, an interferometer capable of measuring the diameter of stars beyond the Sun. Figure
reproduced by permission of the AAS.
The unexpected success of ‘intensity interferometry’ would inspire a host of new projects.
The basic principle behind the intensity interferometer was laid out in Hanbury Brown and
Twiss (1956a), and describes how correlations of intensities (not electric fields) can be used
to measure stellar diameters. First results were reported soon thereafter (Hanbury Brown and
Twiss 1956b), leading to the development of the Narrabri intensity interferometer. With a
188 m longest baseline and blue-sensitivity, this project had a profound and lasting impact on
the field of optical interferometry, measuring dozens of hot-star diameters (e.g. Hanbury Brown
et al (1967a,b, 1970, 1974a), Davis et al (1970)). The small bandwidths attainable with intensity interferometry limited the technique to the brightest stars, and pushed the development
of so-called ‘direct detection’ schemes, where the light is combined before detection to allow
large observing bandwidths. This group would go on to develop the SUSI interferometer.
Dr Charles Townes, inventor of the maser, began a novel interferometer project during
this same time period at University of California at Berkeley. He used heterodyne receivers
as in radio interferometry, but the local oscillators were CO2 lasers operating at frequencies
of ∼27 Hz (or ∼10 µm wavelength), orders of magnitude higher than radio or microwave
oscillators. First experiments were performed using the twin McMath auxiliary telescopes
(separation 5.5 m) at Kitt Peak, AZ; first fringes were obtained on the limb of Mercury in 1974
(Johnson et al 1974) and on stars in 1976 (Sutton et al 1977, 1978, 1979, 1982) where cool
dust shells were detected around many late-type stars. Heterodyne detection also suffers from
bandwidth limitations (like intensity interferometry) as well as additional noise contribution
from laser shot-noise, which becomes progressively worse at higher frequencies. The Townes
group went on to develop the ISI interferometer on Mt Wilson.
Other mid-IR efforts also are worthy of note. An independent project at Arizona using
a kind of aperture masking on a single large telescope (using direct detection) took place
almost simultaneously with the Townes’ experiments (McCarthy and Low 1975, McCarthy
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et al 1977, 1978). In France, Jean Gay and collaborators pursued long-baseline interferometry,
both through heterodyne detection (e.g. Gay and Journet (1973), Assus et al (1979)) and later
direct detection efforts (e.g. Rabbia et al (1990)).
Long-baseline interferometry on a star by directly combining the electric fields before
photon detection (‘direct detection’) was first accomplished in 1974 by Labeyrie (1975),
using a 12 m baseline. This continued the long history of interferometry innovation in France
(starting from Fizeau), and many important experiments have followed. I note that ‘Speckle
Interferometry’ was first described by Labeyrie (1970) and these ideas were also very influential
to the field. However, I will largely limit this review to separate-element, or long-baseline,
interferometry, and will omit comments on speckle. Following this 1974 demonstration in
Nice, the project moved to the Plateau de Calern site and become known as the Interféromètre
à 2 Télescopes (I2T). The I2T made measurements in the visible (e.g. Blazit et al (1977))
and in the near-IR (di Benedetto and Conti 1983, di Benedetto 1985). The Grand I2T (GI2T,
Mourard et al (1994)) began soon thereafter and was developed in parallel with the I2T on the
same plateau, but with larger telescopes (1.5 m) and longer maximum baselines (up to 65 m).
At a time when it was a struggle to simply get two-telescope interferometers working,
considering the array of telescopes needed for imaging was indeed far-fetched. Thus, imaging
using optical interferometry began with aperture masking experiments on large single-aperture
telescopes (in the tradition of Michelson). In aperture masking, a pattern of holes (size r0 in
diameter) is cut in a plate and placed in the pupil plane of a large telescope. The interference
pattern formed thus simulates one from an array of telescopes combined like a Young’s multislit experiment. Baldwin et al (1986) and Haniff et al (1987) showed how aperture masking in
the visible yield data identical to that expected for an imaging array, and produced images of
binary stars using closure-phase imaging. This group, based at the University of Cambridge,
England, would soon begin developing the COAST interferometer, which would succeed in
producing the first image with an aperture synthesis optical array (Baldwin et al 1996). IR
aperture masking at the Keck Telescope (Tuthill et al 2000c) also grew out of work from this
group in collaboration with the U.C. Berkeley ISI team, and exciting unexpected imaging
results from this work have led to much enthusiasm for developing IR imaging capabilities for
long-baseline interferometers such as CHARA and VLTI.
There is one remaining important interferometer lineage to mention, one which led to
the modern development of ‘fringe-tracking’ interferometers such as the NPOI, PTI, and
Keck Interferometers, as well as numerous experimental innovations. The Massachusetts
Institute of Technology and the Naval Research Laboratory built and operated a series of
prototype interferometers, named the Mark I, Mark II, and the Mark III. Shao and Staelin
(1980) reported the first successful active fringe-tracking results, and this group has been most
active at pushing the use of interferometers for precision astrometry. The Mark III was located
on Mt Wilson and was a fully automated interferometer operating in the visible with baselines
up to 20 m (Shao et al 1988). The high efficiency allowed many astronomical programs to
be carried out until it was shut down in about 1993, and is widely considered one of the
most productive interferometers to date. Numerous articles were published covering areas of
astrometry, angular diameters, precision binary orbits, and limb-darkening (e.g. Mozurkewich
et al (1988), Hutter et al (1989), Mozurkewich et al (1991), Armstrong et al (1992), Hummel
et al (1995), Quirrenbach et al (1996)).
I should also mention the prototype interferometer infra-red Michelson array (IRMA),
built at University of Wyoming (Dyck et al 1993). While this instrument did not operate for
very long, those involved were largely responsible for initial success with IR observing at IOTA
and have had lasting impacts at a number of other currently operating US facilities, including
NPOI, PTI, and Keck Interferometers.
Optical interferometry in astronomy
This section was meant to introduce historical interferometers (ones no longer in operation)
which have had a lasting impact on the field, and I have left descriptions of currently operating
interferometers to section 3.6. As discussed in the opening, this review is not meant to document
all the important results from these first generation facilities, but rather to give appropriate
historical background for understanding the current state-of-the-field.
3.2. Overview of interferometer design
Compare Young’s two-slit experiment (figure 1) to what you see in figure 10. We see telescopes
instead of slits and a beam combiner (with relay optics) instead of a screen for viewing the
fringes. In a real interferometer, we must use delay lines to compensate for geometrical delay
introduced by sidereal motion of a star across the sky; in this way, we ‘point’ the interferometer
at the target. In order to successfully interfere light together, each interferometer will have
many subsystems, and in this review we will describe the state-of-the-art developments for the
telescopes, the relay optics, the delay lines, and the beam combination.
Before discussing each of the critical subsystems, the importance of the physical placement
of the telescopes for imaging will be discussed. Many of these issues are discussed in more
detail by Mozurkewich (2000), and here we consider array design from the perspective of
imaging, not for specialized purposes such as nulling or astrometry.
If there were no practical constraints and telescopes could be placed optimally, one could
consider many possibilities. Studies have been published considering distributions based on
optimizing uniformity of (u, v) coverage using three-fold symmetric patterns (used in Keck
aperture masking, Golay (1971)), Reuleaux triangles (used for the Sub-Millimeter Array, Keto
(1997)), a spiral zoom array (considered for the Atacama Large Millimeter Array, see ALMA
memos #216, 260, 283, and 291), and a Y-shaped array (adopted by the VLA). While the first
Figure 10. This schematic illustrates the major subsystems of a modern optical interferometer: the
telescopes, the relay optics, the delay lines, and the beam combination.
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three of these methods offer better Fourier coverage than the Y-shaped array, the ‘imaging’
interferometers of NPOI and CHARA both use a Y-shaped array—why?
For the VLA, an important reason for using a Y-shaped array was a practical one; it was easy
to move the telescopes along railroad tracks in order to cheaply and easily reconfigure the array
geometry. While optical telescopes in arrays do not generally run on tracks (one exception is
IOTA), the desire to transport light to a central facility (see section 3.3.2) leads one to a Y-shaped
geometry where the three-arms of the array are defined by vacuum pipes which relay beams
from the telescopes to the delay lines and combiners. The NPOI interferometer (shown here in
figure 11) has many stations along the three vacuum arms where telescopes can be located, thus
creating a flexible, reconfigurable system capable of pursuing many astronomical programmes.
Another theoretical benefit of this design is that many telescopes can be arranged along each
arm allowing ‘baseline bootstrapping’ for imaging highly resolved targets, a technique where
strong fringes measured between close-by telescopes are used to ‘phase-up’ the fringes on the
longer baselines.
3.3. Critical subsystems technologies
3.3.1. Telescopes. All interferometers need light collectors of some kind. In many cases,
simple ‘siderostats’ are used, whereby a steerable flat mirror directs starlight either directly to
the interferometer or first through a beam-compressor (‘afocal’ telescope). A siderostat has
limited sky coverage and makes polarization measurement difficult (due to the changing,
non-normal reflection angles off the flat), but is thought to offer a more stable structure
for minimizing vibrations and pivot-point drifts for accurate astrometry. More recent
interferometers, such as CHARA, VLTI, and Keck, have chosen traditional altitude-azimuth
(‘alt-az’) telescope designs which give full-sky coverage and potentially salvaging polarization
work. In addition, all interferometer telescopes have incorporated fast ‘tip-tilt’ guiding which
Figure 11. Overhead view of the NPOI interferometer array, to show the Y-shaped array layout,
defined by vacuum pipes extending out to many possible siderostat ‘pads’, or stations. Photograph
reproduced with permission of the Naval Research Laboratory.
Optical interferometry in astronomy
tracks (and corrects) fast jitter of the stellar image, usually using visible-light ‘quad-cell’
detectors. This corrects the first-order term of the wavefront perturbations, aligning the
wavefronts to allow for stable beam combination.
Without high-order adaptive optics, there is little use for a telescope aperture much larger
than the atmospheric coherence length r0 (see section 2.2). Hence, most telescopes in today’s
interferometers are small by ‘modern’ (8 m class) telescope standards. Dedicated visible-light
interferometers (e.g. NPOI, SUSI) have telescope apertures around 12–14 cm in diameter;
near-IR interferometers (e.g. PTI, IOTA, COAST) have apertures diameters around 45 cm.
The recently built CHARA interferometer includes 1 m apertures which can take advantage of
excellent seeing conditions in the IR, and could benefit from adaptive optics correction; the
Keck and VLT auxiliary telescopes were specified to be 1.8 m for similar reasons. However,
interferometry is not just for ‘small’ telescopes anymore, since the world’s largest telescopes,
the two Keck telescopes and also the four VLT telescopes, are now part of the new generation of
optical interferometers. As of 2002, only the Keck Interferometer has observed using adaptive
optics, although the VLT Interferometer will soon possess this capability. See table 3 for a
summary of telescope apertures of today’s interferometers.
3.3.2. Relay optics, delay lines, and metrology. After being collected by the telescopes,
the light must be directed to a central facility for beam combination. While it may seem
trivial to set up a series of mirrors for this purpose, there are many subtle issues that must be
addressed. Traub (1988) discussed how the geometry of the relay optics must not corrupt the
relative polarization of the beams, due to differential phase shifts between the s- and p-wave
reflections from the mirror surfaces for non-normal incidence. One must pay attention to the
issues of mirror and window coatings as well as geometry.
In addition, due to the long path lengths between the telescope and central beam combining
facility, significant differential chromatic dispersion occurs if the light is propagating in air. In
order to combine broad bandwidths, one must either transport the light through a vacuum or
construct a dispersion compensator (Tango 1990, ten Brummelaar 1995), whereby wedges of
glass are inserted into the beam to compensate for air’s index of refraction; a combination of
partial vacuum plus dispersion compensation is also possible. The size of the mirrors in this
optics chain is also important for limiting the effect of diffraction (Horton et al 2001), and
often also sets the field-of-view of the interferometer. Lastly, because of the many reflections,
high reflectivity of the relay optics must be maintained to maximize throughput and sensitivity;
a side-benefit of evacuated relay optics is that the mirrors stay clean.
Because of the Earth’s rotation, the apparent position of an astronomical object is
constantly changing. In order to track this sidereal motion, a movable delay line is needed to
compensate for changing geometrical delay between wavefronts reaching any two telescopes.
The diagram in figure 10 shows this delay line as a right-angle retroreflector, although most
interferometers do not actually use this geometry. The requirements on this system are
amazingly stringent: nanometer-level precision moving at high speeds (>1 cm s−1 ) and over
long distances (>100 m)—a dynamic range of >1010 !
By far the most popular architecture today for the moving delay line is based on the
solution implemented by the Mark III interferometer (Shao et al 1988, Colavita et al 1991).
The retroreflection is produced by focusing the incoming beam to a point coincident with a
small flat mirror (attached to a piezo-electric stack), which reflects and returns a re-collimated
beam; a practical optical system is illustrated in figure 12. This mirror system is mounted on
a flexible stage which can be translated using a voice coil. Lastly, this whole stage is mounted
on a wheeled-cart, which is driven on a rail by linear motors. This system has three nested
feedback loops, driven by laser metrology: precise sub-wavelength control is maintained by
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the piezo-driven small mirror, when this mirror exceeds its normal operating range (∼50 µm)
then offsets are given to the voice-coil stage, and so on. This basic architecture is in use at
most interferometers built in the last 10 years; alternate delay lines include floating dihedrals
mirrors on an air table (IOTA) and moving the beam combination table itself (GI2T).
3.3.3. Beam combination and fringe modulation. Once the beams have been delivered to a
central combination facility and have been properly delayed, there are many ways to actually
do the interference. Here I discuss image-plane and pupil-plane combination (also known
as ‘Fizeau’ and ‘Michelson’ combination, respectively), along with spatial versus temporal
modulation of the fringes themselves.
Figure 13 show these two different ways of detecting fringes in a two-element
interferometer. In one case, an imaging system is used to fill the image plane with the equivalent
Figure 12. Diagram of the most standard delay line architecture used in optical interferometry,
originally from the Mark III interferometer. Figure reproduced from Shao et al (1988, figure 4)
with permission of ESO.
Figure 13. Diagram of image-plane and pupil-plane beam combination techniques. The left
panel shows image-plane, or Fizeau combination, where light from the two telescopes are brought
together in an image plane to interfere, just like Young’s two-slit experiment. The right panel
shows how pupil-plane (or ‘Michelson’) interferometry superimposes the two collimated beams at
a beam-splitter. By modulating the time delay of one beam with respect to the other (e.g. with the
delay line), the interference can be modulated and fringes detected using single-pixel detectors.
Optical interferometry in astronomy
of Young’s fringes. As one moves along the image plane, there is a different relative delay
between the interfering beams, and hence the modulation (fringes). In this scheme, there is no
need to actively modulate the fringes; in fact, atmospheric turbulence will introduce relative
delays and cause the fringe pattern to ‘slide’ back and forth, smearing out the fringes on short
timescales if not stabilized. This combination scheme most closely follows the ‘two-slit’
interferometer analogy developed in earlier sections.
The second scheme, and currently the most common one, is pupil-plane combination,
or ‘Michelson’-style combination. Interestingly, this method is named after Michelson, not
because of his stellar interferometry work (which used ‘Fizeau’ combination, see section 3.1),
but because of the interferometer used in the Michelson–Morley experiment. In this method,
the wavefronts from the two collimated telescope beams are overlapped on a 50/50 beamsplitter.
Depending on the phase relationship of the waves, differing amounts of energy will be
transmitted or reflected at the beamsplitter. Single-pixel detectors can then be used to measure
the energy on both sides of the beamsplitter (the sum of which is conserved). The popular
adoption of this method results largely from the signal-to-noise benefits of using single pixel
detectors. In order to measure the amplitude of the coherence, a dither mirror (often in the
delay line) sweeps through a linear pathlength difference of many wavelengths. The whitelight fringe, or interferogram, can then be recorded, as long as the scanning takes place faster
than an atmospheric coherence time.
When dealing with an array of telescopes, there are more options. ten Brummelaar (1993)
outlined some forward-looking beam combiner designs in the context of the CHARA array
and useful articles by Mozurkewich (2000) and Mariotti et al (1992) also contain extended
discussion on the subject; here I only mention the highlights. The image-plane method can
be extended to arbitrary number of telescopes, as long as the spacings between the beams
are non-redundant, so that each beam-pair will have a unique fringe spatial frequency in the
image-plane. Labeyrie (1996) elaborates on the concept of pupil densification, an idea finding
increasing application in modern interferometry. The pupil-plane method can also be extended,
either by combining the beams ‘pair-wise’ or ‘all-in-one’. In a pairwise-scheme, each telescope
beam is split using beamsplitters and then various combinations are created to measure all the
baselines. In the all-in-one scheme, more than two beams are superimposed and the fringes
from different pairs are distinguished by modulating the delays such that each baseline pair
has a unique fringe temporal frequency in the readout. There are methods which combine
pair-wise with all-in-one and are called ‘partial-pairwise’. Each method has its advantages
and disadvantages, depending on the availability of focal plane arrays, the level of readnoise
versus photon noise, the required calibration precision, etc. However, in general, ‘pair-wise’
detection is the worst method for large number of telescope because the light has to be split
more times (see Buscher (1988), although beware of some important simplifications made in
this analysis).
Coherent beam combination can be discussed more generally depending whether the
interference occurs in the image/pupil plane and whether the telescope beams are co-axial
or multi-axial. I refer the reader to the influential internal ESO report by Mariotti et al
(1992), which explains and defines the useful vocabulary in common use by the European
interferometry community.
We contrast the many imperfect beam combination strategies in the optical with those
adopted in radio interferometry. At radio and microwave frequencies, the signals from each
telescope can be split and re-amplified without introducing additional noise after the initial
coherent detection (radio interferometers do not operate close to the Poisson limit). Hence,
a pair-wise combination scheme can be employed without any loss in SNR. In addition, the
electric field at each telescope can be truly cross-correlated with that from all others leading to
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a kind of Fourier Transform spectroscopy. Further, this can all be done using digital electronics
after fast digitization of the signals. For more information, see the description of the Hat Creek
mm-wave correlator by Urry et al (1985). At the end of this process, the digital correlators can
recover all baselines with arbitrary spectral resolution without lost sensitivity—a dramatically
superior situation than possible in the photon-starved visible and IR regime!
3.3.4. Fringe tracking. An increasingly popular and powerful capability for optical
interferometry is called ‘Fringe Tracking’. To do so, the white-light fringe has to be actively
tracked because atmospheric fluctuations cause the location of the fringe to vary by up to
hundreds of microns on sub-second timescales. There are two levels of tracking these fringes,
one is called ‘coherencing’ and the other is called ‘fringe tracking’, although these terms are
often used rather loosely.
In ‘coherencing’, the interferometer control system will track the interferogram location
to a precision of a few wavelengths. In a scanning interferometer, this will be sufficient to
keep the full interferogram within the scanning range of the delay line. In an image-plane
combiner, this will ensure you are near the peak of the white-light fringe (inside the coherence
envelope set by the spectral bandpass). This can be done on a rather leisurely timescale, since
large optical path distance (OPD) fluctuations tend to occur on slower timescales: update rates
of ∼1 Hz are sufficient except for the worst seeing conditions.
True ‘Fringe Tracking’ (also called ‘co-phasing’) requires tracking OPD fluctuations
within a small fraction of wavelength in RT, and hence requires orders of magnitude faster
response (a timescale which depends on the wavelength and seeing conditions). In the most
common implementation (the ‘ABCD’ method; see Shao and Staelin (1977)), two beams are
combined pairwise while a mirror is stepped at quarter-wavelength intervals. The broadband
white-light fringe is detected at one of the beamsplitter outputs, and fringe data is recorded
synchronous with the dither mirror, resulting in four measurements (A, B, C, D) representing
four different fringe phases. A discrete Fourier Transform (effectively) can be rapidly applied
to the data, resulting in a fringe phase estimate. This offset can be sent to the interferometer
delay line control system to nearly instantaneously correct for atmospheric turbulence (details
in Shao et al (1988), Colavita (1999)). The light from the other beamsplitter output is usually
dispersed and multi-wavelength data is collected. I also refer the reader to Lawson (2000),
where the ABCD method (and other ‘phase’ estimators) or compared to ‘Group Delay’ tracking
methods, which use phase measurements at different wavelengths to measure interferometer
delay offsets.
Historically speaking, active fringe tracking has been important only for the Mark III
interferometers and its successors (NPOI, PTI, Keck Interferometer).
One reason
fringe tracking has not been more widely pursued is because the sensitivity limit of a
fringe tracking interferometer is less than fringe-envelope scanning interferometer. This is
because very short integration times are required to stay on the fringe and hence the source must
be fairly bright; in the fringe envelope scanning method, one has to only keep the interferogram
in the scanning range and thus any given fringe measurement can have a lower SNR. In practice,
this amounts to sensitivity difference of a few magnitudes.
As interferometers become more powerful and seek greater capabilities, fringe tracking
is becoming a standard feature. High spectral resolution interferometry data is possible with
fringe tracking systems, because a broadband white-light fringe can be used for fringe tracking
while the remaining output channels can be dispersed. Normally, this data would have very low
SNR, but if the fringe tracking essentially ‘freezes’ the turbulence, the dispersed fringes can
be detected by integrating on the detector much longer than the typical atmospheric coherence
time. Hence, fringe tracking is a kind of ‘adaptive optics’ for interferometry.
Optical interferometry in astronomy
3.3.5. Detectors. The most desired properties for detectors used in optical interferometry
are low noise and high readout speed, two qualities usually not found at the same time. At the
beginning of optical interferometry, the only visible-light detector was photographic film and
IR detectors were only just invented. Detectors have made incredible advances over the last
few decades, and are operating near their fundamental limit in most wavelength regimes (the
near-IR is a notable exception).
After years of struggling with custom-built photon-counting cameras for visible-light
interferometry work, such as the PAPA camera (Papaliolios et al 1985, Lawson 1994) and
intensified CCDs (e.g. Blazit (1987), Foy (1988)), commercial devices are being sold aimed
at the adaptive optics market which have high quantum efficiencies (>50%), kilohertz frame
times, and read noise of only a few electrons (fast readout CCDs). For beam combination
schemes where single pixel detectors are suitable, Avalanche photo-diodes (APDs) have as
high quantum efficiency as CCDs but can photon-count at rates up to 10 MHz, although the
best commercial devices seem to have an expensive tendency to stop working. This covers
wavelengths from the blue to approximately the silicon cutoff (∼1 µm).
In the near-IR (1–5 µm), there has been amazing progress this decade. After early work
with single-element detectors (e.g. using material InSb), modern interferometers have taken
advantage of technology development at Rockwell in near-IR focal plane arrays made of
HgCdTe, such as the NICMOS3, PICNIC, and HAWAII chips. These arrays have high quantum efficiency (>70%) and can be clocked at ∼MHz pixel rates with as low as 15 e- readnoise.
While not optimal, this represents orders-of-magnitude improvement over photodiodes and
has allowed new kinds of astronomical sources to be observed (most notably, young stellar
objects (YSOs)). The noise can be further reduced by reading each pixel many times, a novel
mode known as ‘non-destructive’ readout. Hence, by reading
√ a pixel n times before resetting,
one can reduce the effective readnoise by approximately n, for n 20. Interferometry
benefits greatly from this capability, since only a few pixels need be readout, allowing large
number of ‘reads’ to be made in a short period of time (interferometry reference Millan-Gabet
et al (1999a)).
Traditionally, the HgCdTe detectors had a cutoff wavelength of 2.5 µm, but recent
molecular beam epitaxy (MBE) processes allow this cutoff to be tuned to much longer (or
shorter) wavelengths (allegedly even beyond 5 µm). For a 2.5 µm cutoff, these detectors must
be operated at liquid nitrogen temperatures (77 K) in order not to be saturated with dark current
from thermally generated electrons. An important recent development is that Raytheon has
begun to compete with Rockwell in this market, and we can hope for even greater advances in
HgCdTe arrays in the coming years as well as possibly even price reductions.
Other materials, such as InSb, can be used for even longer wavelength performance.
At 5 µm and longer wavelengths, thermal background levels are sufficiently high that these
detectors must be readout very rapidly, and can usually work in background-limited mode
(despite >500 e- readnoise). This means that Poisson fluctuations in the thermal background
flux dominate over other sources of noise (e.g. read noise); the only way to reduce the effect
of this background noise is to increase the quantum efficiency of the detector or to reduce the
thermal background load on the detector. Various companies have sold focal plane arrays in the
‘mid-IR’ (∼8–25 µm) over the years, and are not all independent efforts after a complicated
series of company sales (e.g. Hughes, Santa Barbara Research Center, Raytheon, Boeing).
Recently Raytheon has been offering Si : As impurity band conduction (IBC) 320 × 280 focal
plane arrays, which also operate at the background-limit. These detectors must be cooled
below 77 K to avoid high dark currents, and generally use liquid helium. Uniquely, the ISI
interferometer uses a single-element HgCdTe photodiode with high signal level (using CO2
laser local oscillator) which have up to 25% quantum efficiency and a 5 GHz output bandwidth.
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3.3.6. System control. It is not trivial to control all the important subsystems of an
interferometer. Many current interferometers (e.g. ISI, IOTA, PTI, Keck, VLTI) use the VME
RT architecture under the vxworks operating system (Wind Rivers). This allows different
subsystems to be easily synchronized at the millisecond (or better) level. VME systems are
fairly expensive, and some groups (in particular, CHARA) have adopted the RT Linux OS
running on networked personal computers.
3.4. Sensitivity
Optical interferometers are orders-of-magnitude less sensitive than single-dish telescopes.
At visible wavelengths where sensitivity is the worst, current interferometers have a similar
limiting magnitude as the human eye (e.g. V mag ∼6 at NPOI). In this section, we explore the
current and future sensitivities of optical interferometers.
3.4.1. What sets the limiting magnitude? There are three major problems which
limit the sensitivity of today’s interferometers: the atmosphere, optical transmission,
detector/background noise.
The sensitivity is most dramatically limited by the atmosphere which restrict the coherent
aperture size and coherent integration time. We can use the notion of a coherent volume of
photons which can be used for interferometry, with dimensions set entirely by the atmosphere.
The coherent volume has dimensions of r0 × r0 × cτ0 , and hence is very sensitive to the
seeing. Consider average seeing conditions in the visible (r0 ∼ 10 cm, t0 ∼ 10 ms), we can
estimate a limiting magnitude by requiring at least 10 photons to be in this coherent volume.
Assuming a bandwidth of 100 nm, 10 photons (λ ∼ 550 nm) in the above coherent volume
corresponds to a V magnitude of 12.6, which is more than 10 magnitudes brighter than faint
sources observed by today’s 8 m class telescope. Because the atmospheric coherence length
and timescale approximately like λ6/5 for Kolmogorov turbulence, the coherent volume ∝λ18/5 .
Current interferometers cannot achieve this limiting magnitude because of additional
problems. The most important in the visible is low optical throughput due to the large number
of reflections between the telescopes and the final detector. The number of reflections easily
exceeds 10 and is often closer to 20. Even with high quality coatings of 97% reflectivity, we
see that ∼50% of the light would be lost after 20 bounces (0.9720 = 0.54). In practice, current
interferometers have visible-light transmission between 1% and 10%, due the fact that coatings
degrade with time, the need for dichroics and filters with relatively high losses, and some
diffractive losses during beam transport. Of course, detectors also do not have 100% quantum
efficiency. The COAST interferometer has achieved the faintest limiting magnitude in the
visible of ∼9 mag, by optimizing throughput, detector quantum efficiency, and bandwidth (as
another example, the NPOI interferometer which fringe-tracks and uses narrower bandwidths
has a limiting magnitude around ∼6).
Throughput issues can be improved multiple ways. Lawrence Livermore Laboratory
is researching new coatings for mirrors which will have 99% reflectivity at most nearand mid-IR wavelengths. In addition, simplified beam trains with few reflections are being
designed for next generation interferometers. Lastly, the use of fibre and integrated optics could
potentially lead to high throughput systems in the future; these developments are discussed
more fully in section 3.5.
The last major limitation is noise associated with the detection. Some visible light
detectors, such as the photon-counting APDs, are almost perfect in this regard, boasting very
low ‘dark counts’ (<100 ct s−1 ) and high quantum efficiency. However, this is not true in the IR.
Even the best IR detectors have ∼10 e− noise per read. While normal (incoherent) astronomers
Optical interferometry in astronomy
can afford to integrate for minutes or hours to collect photons, interferometrists must readout
pixels within the atmospheric coherence time and thus are strongly limited by readnoise. As
one moves further into the IR (5–10 µm), then thermal background fluctuations dominate the
noise budget. Again, the relatively short coherence times of atmospheric turbulence directly
result in a poor limiting magnitude compared to incoherent detection (i.e. photometry). The
best published near-IR performance of a two-element interferometer was reported by IOTA
(Millan-Gabet et al 1999a) using a NICMOS3 detector: J mag (1.25 µm) 6.9, H mag (1.65 µm)
6.9, and K mag (2.2 µm) 6.2, where J, H are dominated by readnoise and K is dominated by
fluctuations of the thermal background. Soon, these limiting magnitudes will be eclipsed by
the adaptive-optics-corrected Keck and VLT Interferometers which should be able to observe
fainter than 10th magnitude.
There is not much experience yet with mid-IR observations using direct detection. The
ISI heterodyne interferometer has observed stars as faint as ∼360 Jy (LkHα,101 Tuthill et al
(2002)), corresponding to a N band mag of ∼−2.2, limited largely by narrow bandwidths
(λ ∼ 0.002 µm). The VLTI mid-IR instrument MIDI will be capable of broadband
combination and is forecast to have a limiting magnitude of ∼1 Jy (N band mag ∼4) using the
8 m VLT telescopes (assuming the thermal background fluctuations can be well-calibrated for
systematic errors). Shortly before this paper went to press, VLTI reported first fringes with
the MIDI instrument.
A number of new technologies are being explored to push down the limiting magnitude
of optical interferometers, and some of these are described in the next section.
3.5. New technologies and techniques
One exciting aspect to the field of optical interferometry is the aggressive implementation
of new technologies to extend the limits of the sensitivity and calibration precision. In this
section, I will discuss new developments which are impacting optical interferometry.
3.5.1. Spatial filtering and single-mode fibres. The idea to use single-mode fibres in optical
interferometry was originated by Froehly (1982), and work began to implement these ideas in
both France (e.g. Connes et al (1987), Reynaud et al (1992)), and in the United States (Shaklan
and Roddier 1987, Shaklan 1989). Following initial fringe detection in 1991 using the Kitt
Peak McMath telescopes (Coude du Foresto and Ridgway 1992), the FLUOR experiment as
implemented on the IOTA interferometer was a real breakthrough; the amazing improvement
in calibration precision was documented in Coude Du Foresto et al (1997) and Perrin et al
(1998). Currently, the advantages of spatial filtering are being implemented at virtually all
interferometers, and here I briefly explain why it is so important.
Figure 14 shows a schematic of a fibre-based interferometer, as sketched by Coude Du
Foresto et al (1997). When coupling starlight into a single-mode fibre, the coupling efficiency
depends on how coherent the wavefront is from an individual telescope (Shaklan and Roddier
1988). A single mode fibre thus essentially converts phase errors across the telescope pupil
into amplitude fluctuations in the fibre. Once coupled into the single-mode fibre, the light can
be partially split in order to monitor the amount of coupled light (‘photometric’ outputs), and
also can be interfered with light from another fibre using a coupler, the fibre equivalent of a
beamsplitter. The transfer function of the fibre coupler is very stable and not dependent on the
atmosphere; only the input coupling efficiency at each fibre is dependent on the atmosphere.
Hence, the visibility can be measured very precisely (<0.4% uncertainty on V2 reported
by Perrin (2003)) by measuring the fringe amplitude and calibrating with the ‘photometric’
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Figure 14. This figure shows how the FLUOR beam combiner uses spatial filtering and photometric
monitoring to allow precision calibration of fringe visibilities. Figure reproduced from Coude Du
Foresto et al (1997, figure1) with permission of ESO.
This method strongly mitigates the dominant source of calibration error in most optical
interferometers, the changing atmosphere. The atmospheric turbulence must be monitored in
some way when observing with an interferometer, since the coherence between two wavefronts
from two telescopes strongly depends on seeing. However, this is not easy to measure with a
typical interferometer in RT, and hence one must settle for interleaving ‘science’ targets with
‘calibrator’ sources to calibrate seeing drifts during the night. With fibres, the changing seeing
conditions directly cause variations in the fibre coupling efficiencies which are monitored in
RT and corrected for. Figure 15 shows near-IR visibility data on the calibrator star α Boo
using both ‘conventional’ interferometry and the FLUOR fibre optics beam combiner. The
improvement to calibration is indeed dramatic and has had far-reaching effects on the direction
of the whole field of optical interferometry.
There are other ways to implement these calibration advantages than the FLUOR method
shown in figure 14. Monnier (2001) showed how the signal-to-noise can be somewhat improved
by using an asymmetric coupler instead of separate photometric signals. Also, Keen et al (2001)
compared single-mode fibres with spatial filtering by small pinholes in order to determine which
method is superior under different conditions.
I should emphasize that there are many problems and limitations associated with
using single-mode fibres, most notably low coupling efficiencies, high dispersion and poor
polarization stability. Such problems have kept fibre optics from playing an important role for
beam transport (Simohamed and Reynaud 1997), and currently fibres are used only for beam
combining and spatial filtering at specific wavelengths. For instance, silica-based (telecom)
fibres can generally only be used at J (1.25 µm) and H (1.65 µm) bands; the FLUOR experiment
utilized Fluoride glass fibres which can transmit at K band (2.2 µm) and beyond. Advances
in the field of photonic crystals and photonic bandgap materials could lead to new fibres
with low dispersion and high transmission for new interferometry applications, and should be
aggressively pursued.
Optical interferometry in astronomy
Spatial frequency (cycles/arcsec)
Figure 15. (a) This figure shows visibility data for α Boo by the CERGA interferometer () and the
IRMA interferometer ( ), and originally appeared in the Publications of the Astronomical Society
of the Pacific (Copyright 1993, Astronomical Society of the Pacific; Dyck et al (1993), reproduced
with permission of the editors). (b) The incredible gain in calibration using spatial filtering and
photometric monitoring is evident in this figure reproduced from Perrin et al (1998, figure 2(a))
with permission from ESO.
3.5.2. Integrated optics. While combining two telescopes together is straightforward using
fibre optics, it becomes very difficult for multiple telescopes. This is because the light has
been split many times, combined together many times, and the fibre lengths must be precisely
matched and maintained to correct for differential chromatic dispersion and birefringence
An elegant solution to this problem, while maintaining the advantages of spatial filtering,
is the use of integrated optics, the photonics analog to integrated circuits. P Kern and an active
group centred at Grenoble Observatory have pioneered this technique (e.g. Kern et al (1997),
Berger et al (1999), Malbet et al (1999)) and it is finding successful application at a number
of observatories, including IOTA (Berger et al 2001) and VLTI. In these combiners, many
fibres can be mated to a small planar element with miniature waveguides etched in place to
manipulate the light (split, combine, etc). Dozens of beamsplitting and combinations can all
be fit into a few square centimetres—and never needs re-aligned!
While integrated optics can solve the problem of how to combine many beams using guided
optics, it has similar difficulties as fibres of poor transmission, limited wavelength coverage,
dispersion, and birefringence. While the commercial applications for integrated optics in
telecommunications has driven much of the innovation in this field, the astronomy community
must actively engage with the photonics engineers to design custom components which can
overcome the remaining problems for next-generation ‘astronomical-grade’ devices.
3.5.3. Adaptive optics. One critical advance to improve the sensitivity of IR interferometers is
the application of adaptive optics on large aperture telescopes. Generally, visible light photons
are used to measure the wavefront distortions in RT, allowing them to be corrected using a
deformable mirror. Once the aperture is ‘phased-up’, the entire (much larger!) coherent volume
can be used for the IR fringe detection. This method has already been applied on the Keck
Interferometer, where AO systems on the individual 10 m telescopes now allow observations
approaching K mag 10 (and should allow even fainter objects eventually). The major drawback
for this is that there must be a ‘bright’ visible guide star in the isoplanatic patch for the AO
system to use for wavefront sensing, not possible for obscured sources such as YSOs and dusty
evolved stars where the visible source is often too faint (a few AO systems do have IR wavefront
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sensors to mitigate this problem, e.g. Brandner et al (2002)). The maturation of laser guide
star adaptive optics will allow this gain in coherent volume for all IR observing eventually. Of
course, building future interferometers at the most excellent sites (even in space) will be an
increasing priority.
3.5.4. Phase referencing. Phase referencing is a kind of adaptive optics for interferometry,
where a bright reference star is used to measure and correct for atmospheric time delays. This
technique is used in radio interferometry to allow long coherent integrations on targets, via a
fast switching scheme.
Unfortunately, the short atmospheric coherence times make a switching scheme difficult
to implement. A different approach pursued by the Palomar Testbed Interferometer is to use
a ‘dual star module’, where light from two stars are selected and observed simultaneously
using different delay lines. This allows both relative astrometry and phase referencing to be
Very few results have been published on this technique so far, although the technique
is being implemented at the Keck Interferometer and is planned for VLTI. First results from
PTI have been published (Lane and Colavita 1999, Lane et al 2000a, Lane and Colavita
2003), reporting extending the atmospheric coherence time to 250 ms and visibility calibration
precision of 3–7%. Development of this technique will allow very faint limiting magnitudes,
for a small set of sources located within an isoplanatic patch (∼30 arcseconds) of a bright star.
Another method called ‘differential phase’ is being applied soon, where fringes at one
wavelength are basically used to stabilize fringes at all the others. When a source shows
significant wavelength-dependent structure, this technique should prove very powerful. This
is discussed further in the context of extrasolar planet detection in section 5.1.3.
3.5.5. Spectroscopy. Very little has been done in the area of interferometric observations
on spectral lines. The best science results will be reviewed in the next section, however
here I wanted to mention recent developments. Bedding et al (1994) discussed methods for
combining aperture masking with spectroscopy, and the design of the MAPPIT instrument
offers lessons for long-baseline interferometrists. G Weigelt and collaborators have developed
a spectrometer for use on two element interferometers, which allows near-IR molecular
bandheads of CO and H2 O to be spatially probed (e.g. Weigelt et al (2000), Hofmann (2003)).
More interestingly, the AMBER instrument for the VLTI will boast three different spectral
resolutions (up to R ∼ 10 000 across the IR), making observations of individual lines possible
(e.g. Petrov et al (2000)). While the GI2T has had high spectral resolution for years, a number of
other visible-light interferometers, including NPOI and COAST, have modified their combiners
to allow Hα interferometry, following the fascinating results of the GI2T in the 1990s (e.g.
Vakili et al (1998)); see section 4.2.1 for more discussion on this.
3.5.6. New detectors. Single-pixel visible light detectors are nearly ideal in their performance
(e.g. APDs). However, new detectors exist with many of the same advantages of APDs,
but which can also measure the energy of each detected photon (superconducting tunnel
junction detectors, Peacock et al (1997)). Although limited to maximum count rates of
∼10 KHz, current STJ devices offer high quantum efficiency, timing accuracy, and about
∼12% bandwidth energy resolution in the visible and have been used on the sky (Perryman
et al 2001). One obstacle for this technology is that most astronomers want large-format focal
plane arrays with millions of pixels, and present arrays are ∼6 × 6 pixels. These STJ arrays
are small, but large enough to be quite interesting for optical interferometry; this work should
be strongly encouraged.
Optical interferometry in astronomy
For some type of interferometer combiners (e.g. high-resolution spectrometers or
6-telescope imaging), many pixels are needed; unfortunately, APDs and STJs are not
economical for this and CCDs typically still have larger readnoise for fast frame rates. In
this regard, a new development by Marconi may be interesting (Mackay et al 2001). They
have produced a kind of ‘photon-counting’ CCD, which implements on-chip avalanche gain
stages in order to amplify single electrons into large signals. Tubbs et al (2002) report the first
use of these new detectors in astronomy, and the results are promising for interferometry (the
COAST interferometer is currently adapting such a device for their work).
Because of the relatively high readnoise for near-IR detectors, improvements in the next
decade could easily extend the sensitivity of interferometers by a factor of 10. The AOMUX
detector program by Rockwell has just begun, and has the goal of a few electron readnoise at
high frame rates. Keeping pace with these developments will remain a high priority for optical
There are also some developments to create photon-counting near-IR detectors, equivalent
to APDs. Sometimes called solid state photo-multipliers (SSPMs), Eikenberry et al (1996)
described one experiment. Currently, the main drawback with these devices is the low quantum
efficiency, a few per cent. Alternatively, superconducting tunnel junctions can also be used in
the near-IR for photon-counting.
3.5.7. Nulling. Another interferometric technique gaining application is nulling (Bracewell
1978). By introducing an achromatic 180˚ phase shift in one beam, the white-light fringe can
be turned into a white-light null. This has obvious applications for extra-solar planet searches
and zodiacal dust disc characterizations, since removing the bright central star is essential
for detecting faint circumstellar material and companions. The only astronomical results from
nulling have come from aperture masking style experiments (e.g. Hinz et al (1998), Hinz (2001),
Hinz et al (2001b)), and have encouraged aggressive follow-up experiments. In particular, the
Keck Interferometer is pursuing a mid-IR nulling project (Serabyn and Colavita 2001) and
nulling is a central operational mode for the Large Binocular Telescope Interferometer (Hinz
et al 2001a).
3.6. Current and future facilities
In tables 2 and 3, I have summarized all the current and planned facilities (ground-based).
Further discussion of the current field, including documentation of the rising trend of
publications, can be found in Ridgway (2000) where I have found some of the information
for these summary tables. We note that links to all these interferometers can be found on the
well-established ‘Optical Long-Baseline Interferometry News’ website, maintained by Peter
Lawson at NASA-JPL (http://olbin.jpl.nasa.gov).
Each of the currently operating interferometers have unique capabilities and achievements
of note. The GI2T and ISI interferometers are the longest operating interferometers, both
beginning work in the 1980s; notably, the GI2T has uniquely pursued observing of Hα
emission (and remains the only direct detection interferometer with general high spectral
resolution capabilities) and the ISI is the only (published) mid-IR interferometer. The COAST
and NPOI interferometers are currently best optimized for imaging, having incorporated 5 and
6-telescopes, respectively, into their arrays. IOTA is noted for ground-breaking fibre optics and
detector development in the IR. NPOI and PTI have incorporated elaborate internal metrology
to enable ambitious astrometry goals. SUSI has the capability of 640 m baselines and is one of
the only interferometers in the southern hemisphere. Strong progress from the MIRA-I array
marks Japan’s recent efforts in long-baseline interferometry.
Table 2. Current and future optical interferometers: basics (∗ indicates ‘in planning’).
Full name
Lead institution(s)
Center for High Angular Resolution Astronomy
Cambridge Optical Aperture Synthesis Telescope
Grand Interféromètre à 2 Télescopes
Infrared-Optical Telescope Array
Mt Wilson, CA, USA
Cambridge, England
Plateau de Calern, France
Mt Hopkins, AZ, USA
Infrared Spatial Interferometer
Keck Interferometer (Keck-I to Keck-II)
Mitake Infrared Array
Navy Prototype Optical Interferometer
Mt Wilson, CA, USA
Mauna Kea, HI, USA
Mitaka Campus, Tokyo, Japan
Flagstaff, AZ, USA
Palomar Testbed Interferometer
Sydney University Stellar Interferometer
VLT Interferometer (Unit Telescopes)
Georgia State University
Cambridge University
Observatoire Cote D’Azur
Smithsonian Astrophysical Observatory,
University of Massachusetts (Amherst)
University of California at Berkeley
National Astronomical Observatory, Japan
Naval Research Laboratory,
US Naval Observatory
Sydney University
European Southern Observatory
Mt Palomar, CA, USA
Narrabri, Australia
Paranal, Chile
Keck Auxiliary Telescope Array
Large Binocular Telescope Interferometer
Magdalena Ridge Observatory
Optical Hawaiian Array for Nanoradian Astronomy
VLT Interferometer (Auxiliary Telescopes)
LBT Consortium
Consortium of New Mexico Institutions,
Cambridge University
Consortium (mostly French Institutions,
Mauna Kea Observatories, others)
European Southern Observatory
Mauna Kea, HI, USA
Mt Graham, AZ, USA
Magdalena Ridge, NM, USA
Mauna Kea, HI, USA
Paranal, Chile
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Optical interferometry in astronomy
Table 3.
Current and future optical interferometers: capabilities (∗ indicates planned
Size (m)
baseline (m)
47 (100∗)
85 (>100∗)
64 (>250∗)
64 (640∗)
Visible∗, near-IR
Visible and near-IR
Visible, near-IR
Visible, near-IR, 4 µm
Near-IR, mid-IR∗
Near-IR, mid-IR
Near-IR, mid-IR
Visible, near-IR
Near-IR, mid-IR
Recent developments include new IR and visible combiners for the IOTA (first integrated
optics success with stars Berger et al (2001)) and GI2T interferometers, third telescope upgrade
projects for the ISI and IOTA interferometers, 6-telescope operation by NPOI and 5-telescopes
for COAST, and first fringes from the Keck, VLTI, MIRA, and CHARA interferometers. SUSI
has also commissioned a new ‘red’ table, allowing packet-scanning interferometry using APDs.
The FLUOR combiner, so successfully used on the IOTA interferometer, has been moved to
CHARA, and we can expect excellent results soon to take advantage of the greater resolution
and sensitivity.
Indeed, it has been a busy decade for construction and implementation. It is apparent in
table 3 that the current and next generation interferometers boast significantly larger and more
numerous telescope apertures and baselines, and promise to deliver significant new results. In
the area of imaging, CHARA and NPOI will have 6-telescopes spread over hundreds of meters,
to allow imaging capabilities at milli-arcsecond resolution. The VLTI and Keck Interferometers
will have 100 m baselines with adaptive optics corrected primary mirrors, allowing many
new kinds of science to be pursued. In particular, we can expect the first extragalactic
sources, bright AGN and quasars, to be measured at near-IR wavelengths very soon
(probably before this paper goes to press). These new developments are further discussed in
section 5.
Lastly, I will mention recent progress on the next generation of interferometers. The
OHANA project has carried out initial experiments to couple light from Mauna Kea telescopes
into single-mode fibres, the first step in a plan to link the giant telescopes of Hawaii into
a powerful optical interferometer. Major construction for the Large Binocular Telescope
(and Interferometer) has been progressing for many years and is in an advanced stage now.
Importantly, the final design plans for the Magdalena Ridge Observatory are shaping up and
site work for the ∼10 telescope optical array is expected to begin soon. You can find more
information on these ambitious projects in the interferometer summary tables as well.
The next section will review the currently exciting results from optical interferometry, and
give some indication of how the new facilities will impact many areas of astrophysics.
J D Monnier
4. Summary of major scientific results
This section is divided up into two major areas: astrophysics of stars and of circumstellar
environments. Optical interferometers have made substantial contributions in each, and I will
outline recent progress.
4.1. Stellar astrophysics
Optical interferometry has made the greatest impact in the area stellar astrophysics, in particular
the study of nearby single stars. This is not surprising, given the limited nature of singlebaseline interferometers and the limited sensitivity of first-generation instruments. In the last
decades of work, an impressive diversity of investigations have been carried out and here we
document the most successful work.
4.1.1. Stellar diameters and effective temperatures. One of the earliest identified applications
for optical interferometry was directly measuring the effective temperature scale of stars. The
effective temperature is defined such that
L = 4πσ R 2 Teff
where σ is the Stephan–Boltzman constant, R is the radius of the star, and L is the total
bolometric luminosity. Hence, by measuring the angular size of a star and the apparent
luminosity, the effective temperature can be directly calculated. The above equation is often
rearranged in terms of directly observable quantities (independent of distance estimate):
Fbol 1/4
Teff = 2341
where Fbol is the total bolometric flux (in 10−8 ergs cm−2 s−1 ) and θR is the angular diameter
in milli-arcseconds. Empirical calibration of the effective temperature as a function of
spectral type is important since Teff is considered a fundamental parameter of a star, appearing
throughout stellar astrophysics most notably on the Hertzsprung–Russell diagram.
The survey of stellar diameters using intensity interferometry by Hanbury Brown et al
(1974a) still serves as the best resource for the effective temperature scale of hot main sequence
stars. The technique of lunar occultations has traditionally been the other main method for
high-resolution measurements of stellar sizes, as represented by the classic paper by Ridgway
et al (1980).
There are now more than a hundred interferometer diameter measurements, and this
progress is marked in figure 16. Here we see one of the first major results from Michelson
interferometry from the I2T/CERGA interferometer (di Benedetto and Rabbia 1987), the
effective temperature scale of giants. Next to it, is a more recent version of the same diagram
showing the increase in the number of diameter measurements, compiled by van Belle et al
(1999). The effective temperature scale for late-type stars is now well-established, and available
diameter data has been generated by many interferometers (e.g. Mozurkewich et al (1991),
Dyck et al (1996), Perrin et al (1998), Nordgren et al (1999)). A recent cross-comparison
found the data sets from different groups to be statistically consistent (Nordgren et al 2001);
but for the latest spectral types, the visible photosphere is affected by TiO absorption and it is
believed that IR sizes are more representative of the ‘true’ photospheric extent (e.g. Dyck and
Nordgren (2002)).
While giant stars have made easy targets for interferometers due to their large angular
sizes and high luminosities, the census of lower-mass dwarf stars and hotter main sequence
Optical interferometry in astronomy
Figure 16. (a) This figure shows one of the first major Michelson interferometer results (by the
CERGA/I2T interferometer), the effective temperature relations for giants by di Benedetto and
Rabbia (1987, see their figure 1), reproduced here with permission of ESO. (b) Here is reproduced
figure 2(a) from van Belle et al (1999), with permission from the AAS, showing the huge increase
in the number of diameter measurements that now can be used for empirically determining the
effective temperature scale of giants. The x-axis label ‘V–K’ refers to the brightness of the star at
V-band (λ0 = 0.55 µm) compared to K-band (λ0 = 2.2 µm); redder stars have larger V–K colour.
stars remains incomplete. The PTI interferometer recently made first contributions to the study
of K- and M-dwarfs by measuring the diameters of five such stars (Lane et al 2001), although
greater precision (1% diameter errors) is needed to stringently test theoretical models. The
first scientific result from the VLTI interferometer recently contributed to this precious, limited
data set of M-dwarf diameters (Segransan et al 2003), also finding sizes consistent with theory
but lacking precision. To date, there are no published diameters with baselines as long as
the Narrabri intensity interferometer (188m), although the CHARA interferometer did record
fringes on a 330 m baseline in 2001. The new long-baseline capabilities of CHARA, NPOI,
and SUSI should allow progress in some of these areas in the near future.
4.1.2. Limb-darkening, atmospheric structure. Measuring a so-called Uniform Disk (UD)
diameter only requires a single visibility data point for an isolated star, by fitting a oneparameter model. However, stellar photospheres are known to be limb-darkened due to
optical depth effects, and thus UD diameters must be corrected to yield the correct physical
photospheric size. The first attempts to directly measure this was done using the intensity
interferometer on the A1V star Sirius A (Hanbury Brown et al 1974b), however the results
suffered from large errors and were not very definitive. Other attempts have been made based on
looking for wavelength-dependent angular diameters, a sign of limb-darkening (e.g. Ridgway
et al (1982), Mozurkewich et al (1991)).
The visibility curve of a UD star is related to the first Bessel function, and contains an
ever decreasing series of lobes, separated by nulls, as one observes with increasing angular
resolution (see figure 15 for a plot of the first two lobes of this curve). The main difficulty for
limb-darkening studies is that the first ‘lobe’ of the visibility pattern for a star is insensitive
to limb-darkening effects (mathematically, it probes only the 2nd moment of the brightness
distribution; see Lachaume (2003)); measurements beyond for the first null must be made
to unambiguously detect limb-darkening effects (see figure 17 for representative visibility
curves). However, the fringe contrasts at high spatial resolution are necessarily low, and thus it
has been difficult to measure these effects. Precise measurements of Arcturus (α Boo, K1III)
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Figure 17. This figure shows recent results from the NPOI interferometer, which has been optimized
to investigate effects of limb-darkening on stellar profiles. A portion of figure 3 from Wittkowski
et al (2001) has been reproduced here, with permission from EDP Sciences, showing that models
accurately predict the amount of limb-darkening observed in this K5 star.
were made by Quirrenbach et al (1996) using the Mark III employing a novel phase-referencing
method to increase the signal-to-noise near the visibility null; they found reasonable agreement
with model expectations.
The NPOI interferometer succeeded the Mark III, and expanded the wavelength phasereferencing techniques (using strong fringes at one wavelength to allow coherent integrations
on weak fringes). Hajian et al (1998) used these advantages first for limb-darkening studies,
followed more recently by Wittkowski et al (2001). The latter paper presented data with spatial
frequencies sampled well before and after the null and these exemplary results are reprinted
here in figure 17. Again, the atmospheric models were found to be in reasonable agreement
with the interferometry results. A new generation of precision tests of stellar atmospheres
are now being pursued using detailed radiative transfer modelling coupled with thoughtful
interferometer measurements at specific wavelengths and with specific baseline coverage (e.g.
Aufdenberg et al (2002)).
Until recently, most interferometers had only 2-elements and had difficulty to directly
measure the two-dimensional shape of stellar photospheres. In most cases, one assumes
the star is circularly symmetric in order to interpret visibility data taken at differently
projected baselines. The Palomar Testbed Interferometer made the serendipitous, although
not unexpected in retrospect (Hanbury Brown et al 1967b), discovery that the rapidly rotating
star Altair is not circular. Figure 18 shows the oblate spheroid model of this source developed by
van Belle et al (2001). The oblateness is caused by centrifugal ‘bulging’ along the equator and
these measurements offer an independent measure of the projected stellar rotational velocity
v sin i. While red giants and supergiants had been known to deviate from circular symmetric
(see section 4.1.5), this is the first main sequence star found to be non-circular; future ‘imaging’
work should allow new probes of other rotational effects, such as gravity-darkening.
In addition to standard limb-darkening profiles, one has to be concerned with the effects of
molecular lines formed in the photospheres of cool giants (mostly M spectral types), especially
for the Mira variables. It has been known for quite some time (first from speckle interferometry)
that evolved stars appear much larger when observed in narrow spectral channels coincident
with deep TiO bands in the visible regime (e.g. Labeyrie et al (1977)). The origin of this
extension is obvious: at the observing wavelength, the optical depth unity surface is at greater
distance from the star, and hence the apparent size is noticeably larger. Quirrenbach et al
(1993b) made the first systematic study of this effect for late-type stars covering a range of
spectral types, and these results are shown here in figure 19(a).
Optical interferometry in astronomy
Relative Declination (10 m)
Relative Right Ascension (10 m)
Figure 18. Oblate spheroid model for the photosphere of Altair, based on PTI data (from figure 6
of van Belle et al (2001)), reproduced here with permission of the AAS.
Figure 19. (a) This figure shows the Mark III compendium of results measuring photospheric
diameters in and out of strong TiO bands (reproduced from figure 2 of Quirrenbach et al (1993b),
with permission of the AAS). Redder stars (later spectral types) show greater atmospheric extensions
in TiO. (b) This figure shows the most recent data illustrating that Mira stars have strongly
wavelength-dependent diameters (reproduced from figure 2 of Mennesson et al (2002), with
permission of the AAS). The two curves were fit to visibility data taken at 2.2 µm and at ∼4 µm,
and show a greater than factor of two increase in angular size between these two wavelengths.
A more recent, and unexpected, discovery in this vein is that molecules with transitions in
the near-IR are causing large increases in the apparent sizes also. The effect of (most likely)
unappreciated H2 O lines in the coolest M-stars was uncovered by many groups using different
approaches at about the same time. Perrin et al (1999) detected hints of these effects, finding
puzzling deviations from a UD for the O-rich Mira R Leo. Tuthill et al (2000a) found that
R Aqr, another O-rich Mira, was dramatically larger at 3.1 µm than at shorter wavelengths
(an effect seen also in other O-rich Miras, Tuthill et al (1999c)). First results with an L
band (3.75 µm) combiner at IOTA also found a large diameter increase compared to shorter
wavelengths (Mennesson et al 1999).
J D Monnier
A possible explanation for this effect was separately noted by researchers analysing data
from the Infrared Space Observatory (ISO), finding new water features in this part of the
spectrum (e.g. Tsuji et al (1997), Matsuura et al (2002)). Another recent analysis (Jacob
and Scholz 2002) coupled a dynamical model to a simple radiative transfer model and found
complex (time-variable) visibility curves due to molecular effects in the near-IR. Mennesson
et al (2002) have collected data from multiple IR bands (e.g. Chagnon et al (2002)), arguing
the presence of ‘extended gaseous layers’ around O-rich miras; see the dramatic difference in
near-IR sizes observed for R Aqr in the right panel of figure 19.
Thompson et al (2002) have expanded these studies by measuring the sizes of O-rich and
C-rich miras using narrow spectral channels (λ ∼ 0.1 µm) from 2.0–2.4 µm. While they
report O-rich miras are larger near the edges of the band, a different behaviour is observed
for C-rich stars; this tentatively confirms the role of O-bearing molecules (e.g. H2 O). Moving
further into the IR, Weiner et al (2000) (following earlier work by Bester et al (1996)) actually
find stars to be larger in true 11.15 µm continuum channels (using the ISI interferometer)
than at 2.2 µm, a somewhat confusing result since the near-IR wavelengths are expected to
be significantly contaminated by molecular effects. I am aware of even more results which
have not made it to press yet (e.g. adaptive optics at the Subaru telescope, narrow-band IOTA
interferometry), and anticipate rapid progress in this area over the coming years.
4.1.3. Pulsating stellar atmospheres. As discussed earlier, the ‘continuum’ photospheric
size is important for calculating a meaningful effective temperature scale. Another important
consequence of angular diameter work is specific to variable stars: the average physical size
should reveal whether a star is pulsating in the fundamental or first-overtone mode. Distance
estimates have been combined with interferometry data to estimate physical diameters; these
studies typically found ‘large’ sizes consistent with first-overtone pulsation in most cases (e.g.
Haniff et al (1995), van Leeuwen et al (1997), Whitelock and Feast (2000)), although some
sources were found to be fundamental pulsators. This is at odds with both non-linear pulsation
models of Miras (Bessell et al 1996) as well as the persuasive study of variables in the large
Magellanic Cloud by Wood and Sebo (1996). If visible and near-IR diameters are indeed
contaminated by molecular absorption as indicated by recent interferometric results discussed
in the last section, it is possible that the true continuum diameters are small enough to be
consistent with fundamental mode pulsation. The pulsation mode question of Miras has been
debated and ‘settled’ many times, and still more work is needed for a definitive answer.
Pulsating stars, especially the Mira variables (period ∼1 year), are also expected to have
large changes in the photospheric diameters as the luminosity varies (e.g. Bessell et al (1989),
Ya’Ari and Tuchman (1996)). However, because of the difficulty in obtaining a uniform data
set over many years, it is only recently that good pulsation curves of diameters have become
While first hints of phase-dependent diameter changes were reported based on a statistical
analysis of IOTA data (van Belle et al 1996), the first definitive detection of diameter pulsation
came from the COAST group for the O-rich mira R Leo (Burns et al 1998), where a 35%
change in diameter was reported. In figure 20(a), we show more recent results for χ Cyg (also
from COAST) with better temporal sampling of the pulsation curve (Young et al 2000b). Most
recently, Weiner et al (2003) report the first detection of pulsation at mid-IR wavelengths, in
this case for Mira variable o Ceti.
The most significant recent developments are coming from the PhD dissertation of
Thompson at PTI. The high level of automation of the PTI has allowed systematic observations
of a large number of Miras at all pulsational phases, and these data are presented in Thompson
(2002). A first look at the data has been published, and one result from Thompson et al (2002)
Optical interferometry in astronomy
Figure 20. Pulsation curves for two different Mira Variables (at different wavelengths). (a) This
panel shows the large amplitude pulsation of the S-Mira χ Cyg seen at 905 nm by COAST, although
no obvious pulsations were seen at 1290 nm (see figure 2(a) in Young et al (2000b), reproduced
here with permission of Blackwell Publishing). (b) Here we see a well-sampled 2.2 µm pulsation
curve (both open and filled plot symbols) for the O-rich Mira S Lac, from a recent PTI campaign
(see figure 2 in Thompson et al (2002) reproduced with permission of the AAS).
appears in figure 20(b); the diameter of this star has been measured at many phases of the
pulsation, providing an unprecedented opportunity to test non-linear pulsation models. The
full analysis of the PTI data set will shed light on pulsation characteristics as a function of
spectral type (O-rich, C-rich Miras) and wavelength, and we look forward to more of this work
in the near-future.
4.1.4. Cepheid pulsations calibrate period-luminosity relation. Although phenomenologically related to measurements of pulsating AGB stars, observations of Cepheids are quite
distinct in their scientific goals. As has been discussed in Sasselov and Karovska (1994) and
earlier (e.g. Davis (1976)), optical interferometry will play an important role in independently
calibrating the Cepheid distance scale, a crucial element of the cosmic distance ladder. By measuring the changing diameter of a nearby Cepheid and the coeval radial velocity curve through
a pulsation cycle, the distance can be directly inferred via the Baade–Wesselink method.
A flurry of initial results have appeared from GI2T (Mourard et al 1997), NPOI (Armstrong
et al 2001), IOTA (Kervella et al 2001), and PTI (the first definitive detection of Cepheid pulsation; Lane et al (2000b)). However, most current published reports only weakly detect the
pulsation, and definitive results will require longer baselines and/or much higher SNR fringe
This field is rapidly developing, both observationally and theoretically. PTI workers
recently published a second Cepheid results paper (Lane et al 2002), and these remarkable
data are reprinted here in figure 21. The derived Cepheid distances indeed agree with those
measured from the Hipparcos parallax satellite, but the interferometer results are more precise.
The distance uncertainty is now as (or more) affected by our uncertainty in physics of the stellar
atmospheres; fortunately, Marengo et al (2002) report timely theoretical studies of wavelengthand pulsational-phase-dependent variations in Cepheid limb-darkening and emphasize the need
for careful calibration in order to interpret interferometry data accurately. With new observing
campaigns underway at most interferometers, we can expect a rapid development here to really
pin down the pulsational properties of Cepheids in the visible and near-IR. In the near-future,
we can expect the best calibrated Cepheid distance scale to be interferometric.
4.1.5. Imaging stellar surfaces. As discussed in section 3.1, the Cambridge group began
interferometry research by using (visible-light) aperture masking on the William Herschel
J D Monnier
η Aql
ξ Gem
θUD (mas)
θUD (mas)
Figure 21. These data show the most recent pulsation curves for two Cepheid variables, η Aql and
ζ Gem. The data were taken by the PTI interferometer at 1.65 µm and, when combined with radial
velocity data, result in the most accurate distances to these important primary distance indicators.
This figure reproduced from Lane et al (2002, see figure 1), with permission from the AAS.
telescope (WHT) in the Canary Islands (Baldwin et al 1986, Haniff et al 1987) while developing
the COAST interferometer. Interferometric imaging was performed and early results showed
bright features (strong departure from circular symmetry) on the surface of Betelgeuse (Buscher
et al 1990), confirming some previous reports (e.g. Roddier and Roddier (1983)). No longbaseline (separate-element) interferometer would be able to investigate the nature of these
features for years, and the Cambridge masking group has spent more than a decade since
thoroughly investigating ‘hotspots’ on red supergiants and giants.
Over the last decade, it was shown that asymmetries are common (although not
omnipresent) around red supergiants and giants at visible wavelengths (Wilson et al 1992,
Tuthill et al 1997, 1999a), that these hotspots vary on a timescale of months (Wilson et al
1997), and that the asymmetries become less-pronounced (even disappearing) into the IR
(Young et al 2000a). The first image of a stellar photosphere using the COAST interferometer
showed a featureless Betelgeuse (Burns et al 1997).
The hotspots were originally interpreted quite literally, as ‘hot’ patches on the photosphere
from upwellings of large convective elements (Schwarzschild 1975). However, Young et al
(2000a) introduced a new paradigm which is illustrated in figure 22. If the bright features were
indeed caused by literal hotspots, one would not expect the features to completely disappear
in the near-IR. Here, we see an illustrated model where the photosphere is surrounded by a
molecular blanket (e.g. TiO) which is optically thick in the visible, but not in the IR (1 µm).
Inhomogeneities (possibly caused by large-scale convection) allow visible light to escape out
of opacity holes. The results of Dyck and Nordgren (2002) indirectly support this model, by
showing that even ‘continuum’ visible diameters appear contaminated by TiO for late M-stars.
4.1.6. Binary stars and stellar evolution. Binary stars have been an indispensable tool for
astronomers for centuries. Visual and spectroscopic observations yield reliable mass estimates
and form the bedrock of stellar evolution theory (e.g. Eggen (1967)). The advent of speckle
interferometry and optical long-baseline interferometry has led to a remarkable increase in the
data quality and volume of binaries, including many short-period binaries for the first time
(e.g. McAlister (1985), Hartkopf et al (2001)). This work allows unprecedented tests of stellar
evolution models on a case-by-case basis, through sub-1% precision of stellar parameters.
Optical interferometry in astronomy
Figure 22. This figure shows a summary of the results from Young et al (2000a), the most
comprehensive investigation of the origin of surface hotspots on evolved stars. This synthesis shows
nearly coeval images at different wavelengths of the surface of Betelgeuse using a combination of
WHT aperture masking and the COAST interferometer. Surface structures (‘hotspots’) apparent at
visible wavelengths disappear when imaged in the IR. The bottom-right panel shows the schematic
model offered by these authors, where the hotspots are caused not by literally ‘hot’ patches on the
photospheric surface, but rather are caused by ‘opacity holes’ in the molecular envelope (e.g. TiO)
which allow visible light to escape in a patchy pattern. Figures appear here with permission of
J Young.
In addition to the compelling science, the simple nature of their brightness distribution
have made binary stars prime targets for most optical interferometers. Because of the high
spatial resolution, these binaries tend to have short periods and thus full orbital elements can
be determined by tracking the orbit. Not surprisingly, the first aperture synthesis images
by the COAST and NPOI interferometers were of binary systems, and these first results are
reproduced here in figure 23. With sufficient angular resolution, interferometry yields the
angular diameters of the components in addition to the binary separation vector and flux ratio.
While earlier papers had concentrated on individual systems, Hummel et al (1995)
presented orbits of eight systems with separations between 3 and 10 milli-arcseconds using the
Mark III interferometer. For some of the systems, precise mass and luminosity determinations
allowed testing of stellar evolution models.
In order to rigorously test stellar evolutionary models, the highest precision in parameters
is needed. This requires combining data from multiple instruments and techniques in a global
fit, and concomitant attention must be paid to systematic errors. Ideally, the orbital elements
are fit directly to the visibility data and velocities, as implemented by Hummel et al (1998)
and Boden et al (1999). In figure 24, the results from a study of o Leo by Hummel et al (2001)
is presented. This study stands out because it combines interferometry data from the Mark III,
NPOI, and PTI interferometers, as well as radial velocity data, resulting in mass uncertainties
of only ∼0.5%. Stellar evolution isochrones can be put to a serious test for this system.
While long-baseline interferometers allow very close binaries to be partially resolved
(closest is probably 0.002 arcsecond binary TZ Tri; Koresko et al (1998)) and wide systems
to be characterized with incredible precision, this is not always very important. The most
interesting science lies often in measuring unusual binary systems for the first time, such as the
metal-poor double-lined binary system HD 195987 (Torres et al 2002). Attractive targets for
J D Monnier
Figure 23. Baby pictures: first true aperture synthesis images using long-baseline optical
interferometry. (a) The binary star Capella reconstructed at two epochs using the COAST
interferometer (see figure 2 of Baldwin et al (1996) reproduced with permission of ESO). (b) Six
epochs of Mizar A seen by the NPOI interferometer (see figure 4 of Benson et al (1997) reproduced
with permission of the AAS).
Figure 24. Precision binary parameters can be derived by combining interferometry and
spectroscopy, as shown here for o Leo by Hummel et al (2001). The left panel shows the derived
orbit and data from multiple interferometers and the right panel shows a matching stellar isochrone
along with the effective temperatures and luminosities of the two components (see figures 10 and 11
in Hummel et al (2001) reproduced with permission of the AAS).
next generation of interferometer observations include systems with short-lived components
such as Wolf–Rayet stars or YSOs, since much less is already known about the masses of these
Lastly, the push for high dynamic range imaging of binary stars has obvious implications
for detecting low-mass companions, even extrasolar planets, around nearby stars. This topic
will be discussed further in section 5.1.3
Optical interferometry in astronomy
4.2. Circumstellar environments
Interferometers can also be used to probe the environments around stars, both at visible light
and IR wavelengths. As the sensitivity of facilities increase, lower surface brightness features
can be measured, opening up new avenues of research. Advances in interferometric imaging
are particularly relevant here, since gas and dust around stars might not be distributed uniformly
and may be changing in time. While there has not been true imaging accomplished by longbaseline interferometers in this area yet, we have included some of the unexpected recent results
from Keck aperture masking (Tuthill et al 2000c). While equally impressive results have also
appeared using adaptive optics and speckle interferometry (in particular by the Weigelt group),
we highlight the masking results because the observing methods and data reduction closely
parallel that of optical interferometry and more truly reflect future capabilities; they directly
motivate excitement in the potential of interferometric imaging with milli-arcsecond resolution
and point in new scientific directions.
4.2.1. Hα envelopes around hot stars. While almost all the early visible interferometers
focused on angular diameters and binary stars, an interesting exception was observations of
the bright Hα line around Be stars. This emission was expected to be more extended, and
thus more easily resolvable, than the tiny photosphere itself. The envelope of γ Cas was first
resolved by Thom et al (1986) using the I2T, and Mourard et al (1989) saw evidence for an
envelope in rotation by inspecting multiple spectral channels across the line itself using the
GI2T. With a good range of baselines, the Mark III was able to detect definite asymmetries in
γ Cas and ζ Tau (Quirrenbach et al 1993a, 1994); for the latter, MEM was used to visualize the
data as a ‘phase-less’ image and this result is shown in figure 25(a). This data lacked Fourier
phase information, and thus cannot qualify as a true aperture synthesis image; nonetheless,
the maximum entropy procedure provided an innovative and useful tool for visualizing this
asymmetric envelope.
Figure 25. Spectral line observations in Hα around Be stars have found extended, asymmetric
envelopes; the full potential of these kind of investigations remain untapped. (a) Here we reproduce
the ‘image’ of the Hα envelope of ζ Tau from Mark III data (figure 3 from Quirrenbach et al (1994)).
(b) These schematic illustrations of the same envelope at two later dates were based on data from
the GI2T interferometer (see figure 4 from Vakili et al (1998)). Both figures are reproduced with
permission of ESO.
J D Monnier
The high spectral resolution of GI2T later also uncovered asymmetric emission in these
Be star envelopes in Hα (Stee et al 1995), and observed in other lines too (Stee et al 1998).
Figure 25(b) shows results from Vakili et al (1998) (see also Berio et al (1999)) which proposed
that the emission line region is very one-sided and time-variable. The asymmetry in this figure
was derived from the behaviour of the Fourier phases and amplitudes across the line profile;
this amounts to a kind of phase-referencing where the continuum emission surrounding the Hα
emission is used as a reference signal allowing the Fourier phases to be recovered. The origin of
this ‘one-armed oscillation’ could result from radiative effects, the presence of a companion, or
dynamical behaviour in a non-spherical potential; these interesting results should be confirmed
and explored by the current generation of imaging interferometers.
4.2.2. Accretion discs and YSOs. There has been surprising and rapid progress in studies
of how dense accretion discs evolve around pre-main sequence stars. Just five years ago,
simple accretion scenarios incorporating passively-heated flared discs (e.g. Hillenbrand et al
(1992), Hartmann et al (1993), Chiang and Goldreich (1997)) were widely accepted, adequate
to explain the spectral energy distributions (SEDs) of most (low-mass) T Tauri stars and
the higher-mass Herbig Ae/Be systems. However, recent observations with higher spatial
resolution suggest a richer set of phenomena, and has produced much excitement. Direct
observational links are now even being made connecting the fields of star formation and planet
formation, focusing on how accretion discs evolve into protoplanetary discs and finally to
debris discs and planets.
IR interferometry is playing an important role in elucidating the earliest stages of planetary
formation by probing the density and temperature structure of the discs presumably before
planets form. Malbet et al (1998) reported the first resolved YSO, using near-IR data from the
PTI interferometer; the source was FU Ori, a rare type of T Tauri whose emission is dominated
by accretion luminosity, and the disc size was found to be roughly consistent with expectations.
The first ‘normal’ YSO to be resolved with an optical interferometer was the Herbig
Ae/Be star AB Aur, and Millan-Gabet et al (1999b) found the near-IR emission to be much
larger than expected using the IOTA interferometer. These young massive stars have high
luminosities, and thus were the brightest/easiest type of young star to study initially. These
workers published a survey of 15 total Herbig Ae/Be stars which strongly reinforced the initial
finding of large near-IR sizes (Millan-Gabet 1999, Millan-Gabet et al 2001). Also, they found
no evidence for disc structures—intriguingly all the initial data were consistent with spherical
distributions of dust. The much-awaited measurements of classical T Tauri stars would come
soon thereafter from the Palomar Testbed Interferometer (Akeson et al 2000a, 2002). The
few T Tauris measured also showed near-IR emission a few times larger than expected from
accretion disc models. Figure 26 shows the visibility data from these important papers for
AB Aur and T Tau.
While IR imaging of YSOs has only recently become possible with new 3+ telescope arrays
(and nothing published yet), aperture masking on the Keck telescope was able to resolve two
of the brightest examples. Tuthill et al (2001, 2002) imaged the bright source LkHα 101
and found a bright ring of emission, interpreting it as the hot dust at the inner edge of an
accretion disc. While the size was larger than expected for a geometrically thin, optically thick
disc model (e.g. Hillenbrand et al (1992)), the position of the inner edge was consistent with
dust evaporating at temperatures above ∼1500 K when illuminated by direct stellar radiation
(the standard paradigm for dust shells around evolved stars; see Rowan-Robinson and Harris
(1982), Dyck et al (1984)). Further this relation could explain the ‘large’ sizes seen by IOTA
and PTI. Figure 27(a) shows images of LkHα 101 and also of emission-line star MWC 349
by the Keck aperture maskers. We note that a subsequent LkHα 101 paper (Tuthill et al 2002)
Optical interferometry in astronomy
1.61 mas Gaussian
2.62 mas uniform disk
Accretion disk (Ghez)
Accretion disk (Akeson)
Binary companion
Projected baseline (m)
Figure 26. (Left panel): this figure shows the first Herbig Ae/Be star visibility curve ever measured
(see figure 1 by Millan-Gabet et al (1999b)). The IOTA data of AB Aur at 1.6 µm (H-band) and
2.2 µm (K -band) indicated a much larger dust shell or disc than expected from accretion disc
models (dash-dotted curves). (Right panel): the first classical T Tauri stars were observed by the
PTI interferometer and also showed larger than expected IR sizes. Here is the original data for
T Tau itself, an important source but whose analysis is complicated by uncertain calibration due to
the presence of a nearby companion in the interferometer field-of-view (see figure 1(a) by Akeson
et al (2000a)). Both figures are reproduced with permission of the AAS.
Disks around Herbig Ae/Be Stars
Lkha 101
MWC 349A
Ring Inner Radius (AU)
Near-IR Sizes of Herbig Ae/Be & T Tauri Stars
30 AU
-60 -40 -20 0
20 40 60
-60 -40 -20 0
Contours (% of Peak): 1.6 3.1 6.2 12.5 25 50
V380 Ori-A
AB Aur
V594 Cas
T Ori
V1685 Cyg
T Tau N
SU Aur
20 40 60
V1295 Aql
4 AU
Star Luminosity (Lsun)
Figure 27. (a) Aperture masking interferometry was used at the Keck telescope to create these
2.2 µm images of discs around YSOs. The left panel shows the LkHα 101 dust disc with evidence
of a central hole (Tuthill et al 2001), while the right panel reveals the MWC 349A disc viewed
nearly edge-on (Danchi et al 2001). (b) Compendium of measured sizes of Herbig Ae/Be and
T Tauri systems with comparison to theoretical dust sublimation radii (from figure 1 of Monnier and
Millan-Gabet (2002) reproduced with permission of the AAS). Observed sizes are typically many
times larger than expected from ‘classical’ accretion disc models (dotted line). The number of YSO
measurements will rapidly grow with the current size surveys at the Keck and VLT Interferometers.
included the first mid-IR measurements of a YSO disc size (using the ISI interferometer), and
also reported unexpected changes in the LkHα 101 disc emission.
The first results of nulling interferometry observations of Herbig Ae/Be stars (see
section 3.5.7) have been recently published. By masking a single large aperture, Hinz et al
(2001b) found the spatial extent of three sources to be unresolved using a ∼4 m nulling baseline,
an unexpected result. Interestingly, the same simple disc models which under-predicted the
near-IR sizes were shown to over-predict the mid-IR sizes. Clearly, more data is needed to
determine what models are appropriate for discs around YSOs—none of the ‘standard’ ones
J D Monnier
seem to work at near- or mid-IR wavelengths when it comes to spatial observations (the SEDs
can be fit by a number of models).
Recently, Monnier and Millan-Gabet (2002) summarized the current literature of near-IR
disc sizes, and one of the figures appear here as figure 27(b). As indicated above, the most
interesting fact is that nearly all the discs measured are many times larger than expected from
‘standard’ disc models (geometrically thin and optically thick discs), indicating a large dustfree central cavity. A few theorists are incorporating these results into their models, and I
recommend the papers of Natta et al (2001) and Dullemond et al (2001) for further discussion.
Understanding how planetary systems eventually form out of these discs will not be possible
until we are more certain of the initial physical conditions of material within a few AU of
young stars. There is a survey already underway at the Keck Interferometer (and planned soon
for the VLTI) to include many more T Tauri and Herbig Ae/Be stars, and we can look forward
to more developments in this exciting area.
4.2.3. Dust shells and molecules in evolved stars. Long before the dusty discs around young
stars could be observed, interferometry techniques were used to characterize dust shells around
evolved stars. I briefly mentioned the early history of this work in the mid-IR in section 3.1, and
the capabilities of 3 m telescopes were exploited in a series of near-IR speckle measurements
in the 1980s (see especially Dyck et al (1984)).
The most important paper on dust shells was published by Danchi et al (1994), reporting
mid-IR observations of 13 evolved stars by the ISI interferometer. The data were fitted by
radiative transfer models of dust shells, which strongly indicated that mass-loss appeared
continuous around some stars and episodic around others. At the time, mass-loss on the AGB
was still considered to be spherically symmetric and continuous, and these results began to
dislodge this simplified picture.
Subsequent papers concentrated on detailed modelling of individual sources, as the evidence for non-uniform outflows (Monnier et al 1997) and deviations from spherical-symmetry
(Lopez et al 1997) accumulated. Hale et al (1997) found time-variable features in the visibility
curve of the O-rich mira IK Tau, and attributed it to moving dust shells expanding in the outflow.
Figure 28 shows the visibility data and a maximum entropy reconstruction to help visualize the
change in the dust shell morphology; as for the ‘image’ of ζ Tau by Quirrenbach et al (1994),
these radial profiles were reconstructed from phase-less data and so do not necessarily represent a true ‘image’. The changes were consistent with the expected outflow speeds (derived
from maser observations), and allowed an independent distance estimate to this source.
The idea that mass-loss is not uniform in time nor necessarily spherically symmetric found
confirmation using other related techniques. At the same time, near-IR speckle and aperture
masking were being pursued on the new 8 m class telescopes. Imaging of the carbon star
IRC +10216 revealed an incredibly inhomogeneous and asymmetric dust shell (Haniff and
Buscher 1998, Weigelt et al 1998, Tuthill et al 2000c), and enough data has been collected
to see the dust structures evolve with time; I reproduce a figure from Tuthill et al (2000c)
which exploits the diffraction-limit of the world’s largest telescope to image the details of this
nebula (see figure 29). For readers interested in near-IR imaging of dust shells using speckle
interferometry and aperture masking, I refer the reader to many recent papers by the Weigelt
group and the Keck masking team, which will not be reviewed here.
Another recent area of progress is the successful combination of high spectral resolution
with high spatial resolution in the mid-IR. A filterbank spectrometer was constructed and
installed on the ISI interferometer allowing interferometry data to be collected on mid-IR
absorption lines (Monnier et al 2000c,a,b). Polyatomic molecules such as ammonia and silane
form in dense outflows of evolved stars, but at orders of magnitude greater abundance than
Optical interferometry in astronomy
Figure 28. (Left panel): this figure shows temporal changes in the IK Tau visibility curve probing
the dust shell at 11.15 µm as shown by Hale et al (1997, figure 3). (Right panel): these changes can
be visualized through a maximum entropy reconstruction of the dust shell, reproduced here from
figure 4 of Hale et al (1997). Both figures are reproduced with permission of the AAS.
Figure 29. This figure shows temporal evolution of the inhomogeneous and clumpy dust shell
around carbon star IRC +10216 as seen at 2.2 µm using Keck aperture masking (figure 1 from
Tuthill et al (2000b) reproduced with permission of the AAS).
expected (e.g. Keady and Ridgway (1993)). This mystery was partially solved when the ISI
found that the molecules form much further out in the flow than expected, possibly related (for
the case of Silane) to the depletion of SiS onto grains (Bieging and Tafalla 1993). This work
highlights the potential for studying cosmochemistry using interferometry.
J D Monnier
WR 104 at 2.2 µm (Apr 1998)
WR 104 at 2.26µm
Baseline (meters)
Visibility in U-V Plane
V component (105 rad-1)
-40 -20
U component (105 rad-1)
100 -100
Contours (% of Peak): 0.5 1 2 3 4 5 10 30 70
Figure 30. (a) These figures show the 2.2 µm visibility curves, one- and two-dimensional, of
WR 104 observed by Keck aperture masking. (b) This figure shows four epochs of imaging of
WR 104 showing a morphology consistent with an Archimedean spiral rotating at a uniform rate
(from Monnier (1999), figure 13.6)
4.2.4. Colliding winds. While no Wolf–Rayet dust shells have yet to be observed using a
long-baseline interferometer, the discovery of pinwheel nebulae around these stars by Tuthill
et al (1999b) and Monnier et al (1999) using Keck aperture masking represents a new direction
for the current generation of IR imaging interferometers. An example of the spiral pattern seen
in the visibility data of WR 104 is shown in figure 30; from this data (and the closure phases)
images can be reconstructed showing a beautiful spiral structure which appears to rotate on
the sky with a period ∼243 days. These spinning, spiral dust shells result from dust formation
at the interface of two colliding winds in a Wolf–Rayet and O-star binary.
These sources are usually obscured by local dust and thus are faint visible sources
(V 15). The faint V magnitudes make these targets difficult not only for adaptive optics
systems, but also for all current interferometers (except the ISI) which use visible-light detectors
for star tracking and tip-tilt correction. This limitation hopefully will be eliminated at some
facilities by using an IR star-tracker (e.g. IOTA has plans along these lines).
5. The future
5.1. Exciting trends and near-future science potential
As has been said before, now is a practical time for reviewing the major achievements of
optical interferometry, for we are entering a new era boasting facilities with significantly
greater sensitivity, angular resolution, spectral resolution, and wavelength-coverage. In this
section, I will give my views of some of the new capabilities and the expected science returns.
One important trend that must be bolstered is the inclusion of theorists and modellers
in the observations and interpretations of interferometry data. In many areas, the
interferometry observations are outstripping the ready tools for analysis. For example, the
Optical interferometry in astronomy
wavelength-dependent and time-dependent diameters of AGB stars require a combination
of time-dependent hydrodynamical atmospheres and sophisticated radiative transfer codes, a
problem very challenging even with today’s supercomputers. Understanding the hotspots seen
on the surfaces of stars will required three-dimensional simulations of stellar convection.
Accretion disc physics around young stars should include magnetic fields and demand
thoughtful considerations of gas and dust physics in a two- or three-dimensional context. Dust
production in colliding winds is very poorly understood, and poses a formidable numerical
simulation problem. While tackling these difficult physical problems will require the new
high-resolution data from optical interferometers, it is also true that input from the modellers
and theorists is needed to guide and suggest experiments and observing strategies.
Another general comment is that increasing the angular resolution usually means probing
ever decreasing physical scales. Since interferometers often probe scales smaller than an AU,
significant changes in time are expected for even small characteristic velocities (∼km s−1 ). This
poses both a risk and an opportunity: a risk since data must be taken rapidly and efficiently
to accurately capture snapshots of ever-evolving and changing environs, and an opportunity
to include dynamics and time-evolution into our models and understanding. Observing the
dynamics of circumstellar and/or stellar environments allow new physics to be understood,
physics that usually cannot be unambiguously reconstructed from typical data sets. Thus,
I hope that new dynamical information will break theoretical stalemates which paralyse a
number of fields. Interferometers have the opportunity to revolutionize the way we think of
the universe: from distant ‘frozen‘’ images of the past, to a dynamic and engaging unfolding
of the present.
5.1.1. New long-baselines. New long-baselines will allow unprecedented high resolution
measurements on select sources. With sub-milli-arcsecond resolution, one can measure the
diameters of ‘small’ sources which have largely eluded current surveys, such as hot stars and
nearby low-mass stars. Distortions in the photospheric shapes of rapidly rotating stars or binary
stars in nearly Roche-lobe filling systems can be directly detected. Limb-darkening studies of
important objects, such as Cepheids, can be accomplished to put the Cepheid distance scale
of firm direct footing. Further, long-baselines make a variety of exoplanet studies possible,
such as directly detecting 51 Peg b-like planets (‘hot’ Jupiters) or resolving planetary transits
across the stellar disc. The NPOI, CHARA, and SUSI interferometers will possess the longest
baselines in the near-term, while future projects such as the MRO or OHANA might someday
extend the resolution below even 0.10 milli-arcseconds with >1 km baselines.
5.1.2. Imaging. Imaging with optical interferometry is currently tedious at best, and can
only investigate simple objects such as resolved photospheres or binary stars. The 6-telescope
systems of NPOI and CHARA will soon possess the capability of (comparably) excellent
‘snapshot’ coverage, allowing more complicated and higher dynamic range imaging of select
targets. The CHARA array cannot be reconfigured and hence will only image well targets
with the appropriately-sized structures—for a maximum baseline of ∼330 m at 1.65 µm, the
optimum size scale is a few milli-arcseconds. The NPOI interferometer can be reconfigured
to ‘fill-in’ the (u, v)-plane completely over time, and be adjusted for individual sources to
optimally measure the needed visibility and closure phase information. In the longer term, the
auxiliary telescope array at the VLTI and the proposed outrigger telescopes at Keck will allow
even fainter IR targets to be observed. For imaging, the MRO is currently the most ambitious
project in the works, hoping to include >10 telescopes, which would make it the premiere
imaging interferometer in the world.
J D Monnier
Good imaging capabilities would open up new avenues of research, especially in studies of
the circumstellar environments at IR wavelengths. The ability to study discs around young stars
and the time evolution of gaps, rings, or other structures would revolutionize our understanding
of planet formation. At visible wavelengths, imaging spots on the surfaces of other stars is
a major goal, and would allow solar physics to be applied in detail to other stars for the
first time.
The unexpected discoveries of Keck aperture masking justify our optimism that imaging
will uncover many new phenomena that currently are hidden unnoticed in SEDs. For example,
the Wolf–Rayet dust spirals (see figure 30) have only been observed in a few systems, and
represent a new area of study when imaging interferometric arrays are fully commissioned.
However, current imaging work using COAST, NPOI, and IOTA interferometers suffer
from the lack of dedicated software resources. Unfortunately, the decades of software
development in radio interferometry cannot be fully leveraged for optical interferometry,
since radio work now relies largely on phase-referencing techniques not generally available
in the optical. New imaging software is needed which can take into account the unique
nature of optical interferometry data as well as the different nature of our target sources.
The recent adoption of a common data exchange format, defined by the COAST and
NPOI interferometer teams, represent an important first step towards these goals (see
http://www.mrao.cam.ac.uk/ jsy1001/exchange/).
5.1.3. Precision interferometry. This is a rapidly developing area since the advent of singlemode fibres for spatial filtering and ‘dual-star’ phase referencing. When a model of the
astronomical source is well-known, then incredibly precise measurements are possible. The
most potential for this is in the general area of binary stars, where the stars either are pointsources or partially resolved UDs. The case of detecting an exosolar planet is included in this
category, since it can be considered as very high-dynamic-range imaging of a faint companion.
While there are open questions in binary evolution and stellar astrophysics which demand
such high precision, a more popular reason to pursue ‘Precision Interferometry’ is towards
detection of extrasolar planets around nearby stars. There are many ways in which this can be
manifested, and I will outline a few of these.
Narrow-angle astrometry is a comparatively ‘classical’ way to detect an exosolar planet.
Akin to the doppler shift-radial velocity method, precision astrometry attempts to detect the
minute wobble of the parent star as a planet proceeds in its orbit. This can be done by monitoring
the angular distance between a star and a background reference star. In this case, the target
star is normally quite bright and used for phase-referencing to a faint star projected within
an isoplanatic patch from the target (30 arcseconds). Lane et al (2000a) reported the first
measurements of this kind using the PTI (see figure 31). For reference, the motion of Saturn
and Jupiter perturb the Sun ∼1 milli-arcsecond as viewed from 10 pc. This technique will be
applied by the Keck Interferometer and the VLTI interferometer for a planet survey, and there
is talk of pursuing this in Antarctica where the isoplanatic patch is larger and the coherence
times longer (e.g. Lloyd et al (2002), Swain (2002)).
Another method also being aggressively pursued by the Keck and VLTI interferometers is
a multi-wavelength approach to find massive exoplanets by detecting a very slight photocentre
shift between different IR bands due to hypothesized absorption bands in the planet’s
atmosphere (i.e. the differential phase method; e.g. Akeson and Swain (1999), Lopez and
Petrov (2000)). This method has the advantage of using the bright target star as its own phase
reference. However, recent studies of line-of-sight variability of atmospheric water vapour
(Akeson et al 2000b) indicate that differential chromatic dispersion might be more difficult to
calibrate for differential phase methods than originally expected.
Optical interferometry in astronomy
7 Night Residual rms 97 µas
∆ Dec (arcsec)
61 Cyg Primary Size
99 PTI Data
Linear Model
∆µδ = -306 ± 4 ∝as/night
Hipparcos ∆µδ = -315.4 ± 0.4 µ as/night
Residual rms 170 µas
Day Number
Figure 31. State-of-the-art narrow-angle astrometry of the binary 61 Cyg by the PTI. For a period
of one week, the residual astrometric error in declination was ∼100 micro-arcseconds. Figure
printed with permission of SPIE, originally appearing in Lane et al (2000a).
Precision measurements of closure phases can also be used to detect faint companions,
a method which has not received as much attention. As described earlier in this review
(section 2.2.3), the closure phase is formed by summing the interferometer phases on three
baselines around a triangle of telescopes, and this quantity is immune to atmospheric phase
delays. The lack of attention to precision closure phase methods is understandable since
few interferometers possess the requisite minimum of three telescopes. Monnier (2002) and
Segransan (2002) recently discussed how closure phases are immune to dominant calibration
problems of differential phase and that they can also be used to solve for all the parameters of
a binary system without needing to measure any visibility amplitudes. For reference, a typical
closure phase for a binary with brightness ratio of 104 is ∼0.01˚ as long as the component
separation is resolved by the interferometer—the same magnitude effect as for differential
phase methods.
Current published measurement precision of closure phases is only 0.5–5˚ (Benson et al
1997, Tuthill et al 2000c, Young et al 2000a). Improving the three orders of magnitudes
needed to detect even the brightest possible exoplanet is a daunting challenge. While there
are surely unconsidered systematic effects (perhaps due to birefringence or drifts in optical
alignment) which will degrade the sensitivity of the precision closure phase technique, the lack
of any ‘showstopper’ effects, like differential atmospheric dispersion for the differential phase
methods, strongly argues for the further development of the closure phase technique.
5.1.4. Nulling. Another approach being pursued for planet detection is nulling (Bracewell
1978). The initial nulling experiments with the MMT (Hinz et al 1998) have continued (Hinz
2001), and ultimately will be applied on the Large Binocular Telescope Interferometer (Hinz
et al 2001a). This project is still many years away, but offers an alternative approach to the
‘precision’ phase methods above.
In the nearer term, the Keck Interferometer will be applying nulling in the mid-IR in order
to measure and characterize the zodiacal dust around nearby stars. This source of IR radiation
J D Monnier
is expected to be the dominant background for an eventual space-based planet detection
interferometer, the so-called ‘Terrestrial Planet Finder’ (TPF) mission (more information in
section 5.2). Serabyn and Colavita (2001) describe the ‘fully symmetric’ nulling combiner
being implemented on the Keck Interferometer, and initial on-sky tests are expected to begin
in 2003. A more complete description of the observing strategy and expected sensitivity has
been documented in Kuchner and Serabyn (2003).
Nulling can also be applied on large single-apertures, and then are called nulling
coronagraphs. New clever designs in coronagraphy are competing with nulling interferometry
for space mission concepts to detect terrestrial planets around other stars, and I recommend
interesting papers on optimally shaped and apodized pupils (Spergel 2002, Nisenson and
Papaliolios 2001), bandpass-limited image-masks (Kuchner and Traub 2002), and phase-maskbased approaches (e.g. Guyon et al (1999), Rouan et al (2000)).
5.1.5. Spectroscopy. There have been only a few significant results combining spectroscopy
and interferometry; fortunately, this is about to change. The near-IR AMBER instrument,
slated to arrive at the VLTI interferometer in 2003, will combine three telescope beams together
and disperse the light with three different spectral resolutions, the maximum is R 10 000.
This resolution will allow interferometry on individual spectral lines in the 1–2.5 µm regime,
opening up shock-excited emission lines, CO-absorption/emission features, and even emission
from YSO jets to be probed in novel and exciting ways for the first time. We can expect the
value of interferometric observations to be greatly enhanced by these new capabilities.
5.1.6. Polarimetry. Imaging stars in polarized light with interferometers also promise
fascinating new insights into many areas of astrophysics, although this capability is difficult
to implement with current interferometers. Vakili et al (2002) discuss interesting applications
of combining the high spectral resolution of AMBER with polarimetry, and highlight the new
capabilities for imaging scattered light and potentially even measuring stellar magnetic fields
from the Zeeman effect. Experimental efforts (Rousselet-Perraut et al 1997) in this area have
been very limited compared to the theoretical progress (Rousselet-Perraut et al 2000); this
situation should be remedied soon.
5.1.7. New observables. Along with greater spectral coverage and more telescopes come new
interferometric observables. While section 5.1.3 discussed possible applications of differential
phase and differential closure phase, there are other interferometric observables yet to be
exploited for precision interferometry.
Measuring the diameter of a star by precisely locating the first null of the visibility
pattern is immune to amplitude calibration errors. This could be done by using a wellcalibrated spectrograph to search for the null, either measuring fringe amplitudes or looking
for the signature phase-flip across the null (e.g. Mozurkewich, private communication). This
technique is similar to the method of A Michelson in measuring the diameter of Betelgeuse
(Michelson and Pease 1921), where the baseline was adjusted in order to find the visibility
minimum as detected by his eyes.
The closure amplitude (requires sets of 4-telescopes) is an important quantity in radio
interferometry to compensate for unstable amplifier gains and varying antenna efficiencies that
can be linked to individual telescopes (e.g. Readhead et al (1980)). Closure amplitudes are not
practical for current optical interferometers partially because most fringe amplitude variations
are not caused by telescope-specific gain changes but rather by changing coherence (e.g. due to
changing atmosphere). However, the introduction of spatial filtering (e.g. single-mode fibres)
Optical interferometry in astronomy
should make the closure amplitude a useful tool for optical interferometry soon (see discussion
in Monnier (2000)).
Necessarily, most new observables have yet to be used in practice or described in
print. I mention here a few possibilities that this author has considered to encourage future
experimentation. For instance, it may be possible to use closure amplitudes in the case when
fringe jitter causes loss of visibility contrast in a fringe-tracking interferometer, due to the way
in which small random phase errors degrade coherence. Also, the closure differential phase is
a recently defined quantity (Monnier 2002), introduced to overcome one limitation of current
phase-referencing techniques, namely, that differential phase (and differential closure phase)
methods requires assumptions about the source structure of the phase calibrator.
5.1.8. Sensitivity (Keck and VLTI). Another area where we expect immediate progress is in
observing new classes of faint objects for the first time. The Keck and VLTI interferometers will
have the capability of observing sources as faint as K ∼ 11 magnitude (down to K ∼ 20 with
phase referencing), opening up extragalactic sources for the first time. By the time this review
is printed, I expect that the first optical interferometric observations of the core of an AGN
will be announced. Size measurements of AGN should offer new constraints on models of the
IR continuum and, when coupled with high spectral resolution, could determine the physical
origin of observed broad line regions and possibly even measure dynamical black hole masses.
In terms of galactic sources, this increase in sensitivity will allow a broad census of sources
to be taken, including YSOs spanning a broad range of ages, luminosities, and distances and
binary systems of all masses. For instance, IR observations of pre-main-sequence binaries
allow unique probes of the evolution of binary fraction (e.g. Ghez et al (1993)) as well as
important measurements of YSO masses (Tamazian et al 2002). I expect interferometer
observations to play an increasingly important role in this area as the sensitivity increases.
Of course the additional sensitivity will permit new projects too, such as tracking the
motions of stars orbiting the black hole at the centre of the Milky Way with an order of
magnitude greater precision than possible today with single-aperture telescopes (e.g. Schödel
et al (2002)). Precision astrometry may allow even new tests of general relativity near super
massive black holes at the centre of nearby galaxies.
In addition, the MIDI instrument for the VLTI will allow sensitive measurements in
the mid-IR for the first time. While the ISI interferometer pioneered interferometry in this
wavelength range, MIDI+VLTI will be first to probe a wide range of sources with resolution of
∼0.01 arcseconds and down to N ∼ 4 mag (100× fainter than the ISI). Mid-IR observations
are sensitive to emission from relatively cool dust and can peer through thicker layers of
dust than possible in the visible or near-IR. There are great possibilities for advancing our
understanding of young and evolved stars both, and studying dust distributions in a variety of
5.2. Space interferometry
The greatest limitations to optical interferometers arise from atmospheric turbulence. It
dramatically limits the sensitivity, the ability to do imaging, and forces the engineering to
be clumsy and complicated. Space is naturally an ideal place for interferometry, with no
atmosphere to corrupt the phase nor limit the coherent integration time. And long-baselines are
obviously possible by combining light intercepted by separate spacecraft flying in formation.
5.2.1. Critical technologies needed. In order to successfully build space interferometers,
many technologies must first be developed. To this day, there has not been any dedicated space
J D Monnier
interferometer flown (except for the fine guidance sensors on the Hubble space telescopes; e.g.
Franz et al (1991)).
For interferometers deployed on a single structure, one has to contend with truss vibrations,
thermal and gravitational gradients, and an unusually large number of mechanisms (failures
of which could end the mission). There are issues with propellant and power consumption for
maneuvering the array to point around the sky. The Space Interferometry Mission (SIM) is
in advanced planning stages and is being designed to measure accurate positions of stars with
micro-arcsecond resolution. SIM is a ‘simple’ 2-element interferometer on a deployable truss
(∼10 m maximum baseline), and will be the first space mission to attempt space interferometry.
Ultimately, one would want to have baselines much longer than ∼10 m, and this will
require separate, free-flying spacecraft. For a space interferometer consisting of ‘free-flyers’,
there are other problems. For instance, maintaining the physical distances between space
telescopes to sub-micron tolerances is indeed a challenge. Probably this cannot be done;
however by monitoring the spacecraft drifts in RT using laser metrology, the changing distances
can be compensated for by onboard (short) delay lines. Some engineering missions have been
proposed to test ideas, but have yet to really get-off-the-ground (e.g. the NASA Starlight
mission was recently cancelled). NASA and European space agency (ESA) should give such
a test mission a high priority since the science potential for a free-flyer interferometer is so
much greater than for one limited to a single structure.
5.2.2. Review of current NASA and ESA missions. There are a number of mission concepts
involving space interferometry being considered by NASA and the ESA. As mentioned before,
the only one in advanced design stages is the NASA SIM. In table 4, I summarize some of the
missions that are being proposed, and their main science drivers. Considering the unreliability
of expected launch dates, I have omitted these from the table—it is unlikely any of these will
fly before 2010 (2020?).
NASA and ESA have spent much energy on designing missions to detect Earth-like planets
around nearby stars, and to measure their crude reflectance (or emission) spectra. With luck, an
extrasolar planet spectrum could encode distinctive atmospheric spectral features indicating the
presence of life (biomarkers) on the distant planet (e.g. Woolf et al (2002)). While originally
envisioned as an IR interferometer mission, concepts involving a visible-light coronagraph
have been proposed lately. This mission is known at the TPF in NASA, and as IRSI-Darwin at
ESA. The summary table also includes a few TPF follow-on missions, such as ‘Life Finder’.
These missions are very futuristic, and testify to NASA’s ebullient imagination.
Another area of interest is imaging the far-IR and sub-millimeter sky at high angular
resolution using space interferometry. These wavelengths are difficult to access from the
ground due to water absorption in the atmosphere. Because of this, the angular resolution of
current observations are very limited (∼30 arcseconds); compared to all other wavelengths,
the sky has been surveyed with the lowest resolution in the far-IR.
The proposed NASA mission ‘Submillimeter Probe of the Evolution of Cosmic Structure’
(SPECS) would be a separate-telescope space interferometer (possibly tethered together and
not ‘free-flying’) designed to map the sky with great sensitivity at a resolution comparable
to that currently achievable at other wavelengths (∼0.010 arcseconds). This would avoid the
confusion-limited regime encountered by current low-angular-resolution galaxy count surveys,
and allow the evolution of cosmic structure to be investigated back to high redshift. The SPIRIT
mission is meant as a precursor to SPECS to test out various aspects on a single platform.
The x-ray community has also proposed a space interferometer, which would boast microarcsecond resolution and be capable of studying the hot material at the event horizon of
nearby black holes. Bolstered by successful lab experiments (Cash et al 2000), plans for a
Optical interferometry in astronomy
Table 4. Proposed space interferometers.
Full name and primary science drivers
Space Interferometry Mission
Precision astrometry; exosolar planets
Fourier–Kelvin Space Interferometer
Find Jovian planets (nuller); map circumstellar discs
Test free-flying concept for ESA IRSI-Darwin mission
Infra-Red Space Interferometer (one concept: Darwin)
Image terrestrial planets (IR nuller); measure spectra
Terrestrial Planet Finder
Image terrestrial planets (IR nuller); measure spectra
Space Infrared Interferometry Trailblazer
Far-IR, sub-mm galaxy counts; precurser to SPECS
Submillimetre Probe of the Evolution of Cosmic Structure
High-resolution map of high-Z universe (far-IR, sub-mm)
Stellar Imager
Image surfaces of stars (visible, ultraviolet)
Micro-Arcsecond X-ray Imaging Mission
Map black hole accretion discs and event horizons (x-rays)
MAXIM Pathfinder
Demonstrate feasibility of x-ray interferometry; achieve 100 µ-arcsecond resolution
Life Finder
Search for biomarkers in planet spectra; TPF extension
Planet Imager
Image surfaces of terrestrial planets, 25 × 25 pixels
(requires 6000 km baselines, futuristic!)
IRSI-Darwin (ESA)
MAXIM Pathfinder
free-flying x-ray interferometer called the Micro-Arcsecond X-ray Imaging Mission (MAXIM)
have begun. Controlling distances between macroscopic mirrors to picometre-precision, as is
needed for x-ray interferometry, is indeed a daunting challenge. However, a MAXIM precursor
mission with only a few metre baseline would have orders of magnitude greater resolution than
the Chandra x-ray telescope and stands some chance of being flown.
5.3. Future ground-based interferometers
While it is interesting to speculate about the future of space interferometry, we recognize that
it will be expensive, difficult, and slow-paced. In the next 10 or 20 years, we can expect
more affordable and rapid progress to be possible from the ground. In this concluding section,
I review some of the necessary characteristics of an optical VLA (OVLA). Ridgway (2000)
discusses many of these considerations, and I refer the reader to his interesting report for further
5.3.1. Design goals. The main design goal of a next-generation optical interferometer array
will be to allow the ordinary astronomer to observe a wide-range of targets without requiring
extensive expert knowledge in interferometer observations. An imaging interferometer with
great sensitivity could fulfill this promise by providing finished images, the most intuitive data
format currently in use. It will not be a specialty instrument with narrow science drivers, but a
general purpose facility to advance our understanding in a wide range of astrophysical areas.
J D Monnier
5.3.2. OVLA. One way to achieve this design goal is to scale up the existing arrays. Simply
put, this main goal will require an array with a large number of telescopes (20 to allow
reliable aperture synthesis imaging) and with large-aperture telescopes corrected by adaptive
optics (preferably using laser guide stars for full-sky coverage), allowing a reasonably faint
limiting magnitude (roughly speaking, brighter than ∼15th magnitude in the IR with no phase
This array would likely be reconfigurable, like the radio VLA, to allow different angular
resolutions to be investigated. The longest baselines should cover a few kilometres (∼0.1 milliarcsecond resolution in the near-IR). The main limitation of such a system will be a small fieldof-view, typically limited to the diffraction-limited beam of an individual telescope (for 10 m
class telescopes, the instantaneous field of view would be only about ∼50 milli-arcseconds)—
although mosaicing would be possible, as in the radio. There are schemes which can image a
larger field simultaneously, but are probably not very practical.
With an even larger (billion-dollar) budget, one can partially combine the goals of
interferometry with the community priority for a 30 m diameter telescope. This clever idea was
recently proposed by R Angel and colleagues at the University of Arizona. In their ‘20/20’
scheme, light from two extremely large telescopes (diameter >20 m) would be combined in
a Fizeau combination scheme, patterned after the Large Binocular Telescope, maintaining a
much larger field-of-view (maybe ∼30 arcseconds, limited by atmospheric turbulence) with the
resolution of the two-element interferometer. Further, this scheme maximizes raw collecting
area and would boast potentially incredible sensitivity (>20 mag!). One demanding feature
of this design is that the two 20 + m telescopes would have to smoothly move along a track in
RT to maintain the large field-of-view; this may not be impossible, but is surely an interesting
complication. Further, the imaging advantages of this system only work when the 2-telescope
baseline is 5–10× as large as the telescope diameter, and hence the ‘20/20’ interferometer
would have maximum baselines of only a few hundred metres at most, not much better than
current interferometer arrays. While granting that this system could allow much fainter objects
to be observed, this option would cost many times more than a dedicated OVLA system
described above, and would have much poorer angular resolution.
5.3.3. Technological obstacles needed to be overcome. If optical interferometry is to continue
its impressive growth over the coming decades, important breakthroughs must be made in
critical areas. Here, I briefly list a few obvious improvements which would make an OVLA
more affordable.
The main advance needed to make the OVLA affordable will be the development of ‘cheap’
large-aperture telescopes with adaptive optics. Currently, it costs millions of dollars to build
even a 4 m-class telescope—without adaptive optics. Advances in lightweight mirrors with
adaptive optics designed-in from the beginning may change the economics of the situation.
Another area which could revolutionize optical interferometry is advances in photonic
bandgap fibre materials (e.g. Mueller et al (2002)). These materials offer possibility of
extremely wide-bandwidth, low dispersion and low-loss single-mode fibres, which could open
up the possibility of practical fibre delay lines. Such an advance would greatly simplify
the optical beam-train and engineering of an optical interferometer, making projects such as
OHANA straightforward. This would put optical interferometry on more similar footing as
radio interferometry, where cable delay lines (either coaxial or fibre) are routinely used.
Combining dozens of telescopes may not be practical using bulk optics, and solutions
involving integrated optics should be pursued. The main limitation of this technology is
restricted wavelength coverage, currently only proven shortward of 2.2 µm. Development of
materials (e.g. lithium niobate) and fabrication processes that can extend the coverage into
Optical interferometry in astronomy
the thermal IR (1–5 µm) would mean that a general purpose interferometer could be built
around an integrated optics combiner. Work is currently underway in Europe towards this
end, in particular in pursuit of mid-IR nulling capabilities for the ESA IRSI-Darwin mission
(Haguenauer and others 2003).
Lastly, improved IR detectors are crucial to maximizing the scientific output of a future
interferometer. It has already been discussed here (see section 3.3.5) that near-IR detectors
remain limited by avoidable detector ‘read’ noise, and a future OVLA must have better
6. Conclusion
After decades of development, optical interferometry is now poised to play a major role in
mainstream astronomy. The emergence of well-funded interferometer ‘facilities’, in particular
the very large telescope interferometer and the Keck Interferometer, promise to revolutionize
the impact of high-resolution observations in many areas of astrophysics. Clearly, the main
beneficiaries will be stellar astrophysics and galactic astronomy, in particular the areas of star
and planet formation, fundamental stellar properties, and all stages of stellar evolution. In
addition, we can look forward to the first extragalactic results.
Although the Keck and VLT interferometers hold immense promise, the field is currently
driven forward by the activities of many other smaller groups, and scientific results will be
dominated by these workers for the near-future. While many experimental (astro)physics
fields have matured to the point where future progress rests in ‘big science’ collaborations
and national research centers (e.g. NASA), optical interferometry represents one of the few
healthy and active ‘experimental astrophysics’ endeavors left in astronomy where universitybased groups continue to make important technical innovations and astronomical discoveries.
It is widely acknowledged that astronomy as a whole is experiencing a golden age of
progress, spurred on by observational advances across the electromagnetic spectrum. Optical
interferometry has expanded in response to its own promising initial results, and we in the field
look forward to exploiting the significant infrastructure buildup just now being completed.
I hope that the next review of optical interferometry will vindicate my optimism in the field,
and that the pioneering discoveries reported here presage even grander exploits. It is safe to
predict that the next decade will be critical to the field of high-resolution optical astronomy,
since the scientific impact of current facilities will wholly determine whether the substantial
funding required for an OVLA can be justified to the international astronomical community.
Firstly, I must apologize for omitting many important works due to space constraints, especially
in the areas of speckle interferometry. I thank J-P Berger, R Millan-Gabet, and P Lena for
a careful reading of the manuscript and important suggestions. Also, I acknowledge useful
conversations with E Pedretti, A Boden, P Tuthill, and F Malbet. Lastly, I recognize formative
discussions with S Ridgway, C Haniff, D Buscher, and D Mozurkewich, whose ideas have
helped shape my perspective of the field of optical interferometry, especially on the future of
an OVLA.
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