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Serb. Astron. J. } 181 (2010), 1 - 17 UDC 520.3613DOI:
10.2298/SAJ1081001J Invited review
ASTRONOMICAL OPTICAL INTERFEROMETRY.I. METHODS AND
INSTRUMENTATION
S. Jankov
Astronomical Observatory Belgrade, Volgina 7, 11060 Belgrade 38,
SerbiaEmail: [email protected]
(Received: November 17, 2010; Accepted: November 17, 2010)
SUMMARY: Previous decade has seen an achievement of large
interferometricprojects including 8-10m telescopes and 100m class
baselines. Modern computerand control technology has enabled the
interferometric combination of light fromseparate telescopes also
in the visible and infrared regimes. Imaging with milli-arcsecond
(mas) resolution and astrometry with micro-arcsecond (as)
precisionhave thus become reality. Here, I review the methods and
instrumentation cor-responding to the current state in the field of
astronomical optical interferometry.First, this review summarizes
the development from the pioneering works of Fizeauand Michelson.
Next, the fundamental observables are described, followed by
thediscussion of the basic design principles of modern
interferometers. The basic inter-ferometric techniques such as
speckle and aperture masking interferometry, aperturesynthesis and
nulling interferometry are disscused as well. Using the experience
ofpast and existing facilities to illustrate important points, I
consider particularly thenew generation of large interferometers
that has been recently commissioned (mostnotably, the CHARA, Keck,
VLT and LBT Interferometers). Finally, I discussthe longer-term
future of optical interferometry, including the possibilities of
newlarge-scale ground-based projects and prospects for space
interferometry.
Key words. Instrumentation: interferometers Methods:
observational Tech-niques: high angular resolution
1. INTRODUCTION
An optical (visible and infrared) long-baselineinterferometer is
a device that allows astronomers toachieve a higher angular
resolution than it is possi-ble with conventional telescopes. In
fact, at wave-length , the resolution of a single telescope
withaperture diameter D scales as D/ while the reso-lution of
two-telescope interferometer scales as B/,where the baseline B is
the distance between the tele-scopes. However, more severely, both
are limited bythe atmospheric turbulence (seeing typically 1
arc-second). For this reason, a single aperture telescope
can be used as an interferometric device, but usuallyit is
composed of an array of at least two telescopes,which sample the
wavefronts of light emitted by asource at separate locations, and
redirect starlight toa central location in order to recombine the
sampledwavefronts and to produce interference fringes. Thecontrast
of interference fringes, or visibility, variesaccording to the
characteristics of the light source(for example, the size of a star
or the separation be-tween two stars in a binary system) and
accordingto the length and orientation of the
interferometersbaseline, the line connecting the two telescope
aper-tures. It is possible to take measurements from manydifferent
baselines, most easily by waiting while the
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S. JANKOV
Earth rotates. In addition, most of the new inter-ferometers
have more than two mirrors in the arrayand can move the apertures
along tracks.
The implementation of interferometry in opti-cal astronomy began
more than a century ago withthe work of Fizeau, who outlined basic
concept ofstellar interferometry, how interference of light canbe
used to measure the sizes of stars. He suggestedthat one could
observe stars using a mask with twoholes in front of a large
telescope and that it shouldbe possible to obtain some information
on the angu-lar diameters of these stars (Fizeau 1868). The
firstattempts to apply this technique were carried outsoon
thereafter (Stephan 1874), although the largestreflecting telescope
in existence at the time, the 80cmreflector at the Observatoire de
Marseille, could notresolve any stars with a mask that had two
holes sep-arated by 65cm, and only the upper limit (0.158 arc-sec)
of stellar diameter could be derived. However,with a Fizeau
interferometer on the 1m Yerkes refrac-tor, Michelson (1891a,b)
measured the diameters ofJupiters Galilean satellites. Following
Michelsonssuccess Schwarzschild (1896) managed to resolve anumber
of double stars with grating interferometeron 25cm Munich
Observatory telescope, while notlong afterward, on the 2.5m
telescope on Mt. Wilson,Anderson (1920) determined the angular
separationof spectroscopic binary star Capella ( Aur).
Although the first measurement of a stellar an-gular diameter
was performed on the supergiant starBetelgeuse ( Orionis) with
Michelsons 6m stellarinterferometer in 1920 (Michelson and Pease
1921),the optical interferometry was slowly evolving from
adifficult laboratory experiment to a mainstream ob-servational
technique. Following the success of the6m interferometer, Pease
(with Hale) constructed a15m interferometer but this experiment was
not verysuccessful. Due to the disappointing results from the15m
interferometer, it would be decades before sig-nificant
developments inspired new activity in theoptical field. The real
difficulty is to combine thebeams in phase with each other after
they have tra-versed exactly the same optical path from the
sourcethrough the atmosphere, each telescope, and furtherto the
beam recombination point. This has to bedone to an accuracy of a
few tenths of the wave-length, which in the case of visible light,
is not an ob-vious task, particularly because of atmospheric
tur-bulence which makes the apparent position of a staron the sky
jitter irregularly. This jitter often causesthe beams in different
arms of the interferometer tooverlap imperfectly or not at all at
any given mo-ment. For these reasons, the optical interferome-try
requires extreme mechanical stability, sensitivedetectors with good
time resolution, and at least asimple adaptive optical system to
reduce the effectsof atmospheric turbulence. For these reasons,
onlythe technologies emerging at the end of 20th centuryallowed the
full application of optical interferometryin astronomy.
Meanwhile, advances in radar technique dur-ing World War II
stimulated rapid development ofradio interferometry beginning with
the first radiointerferometer built by Ryle and Vonberg (1946).The
main reason for tremendous success concern-
ing astronomical applications is that the effects ofatmospheric
noise are much less pronounced at ra-dio wavelengths. Paralelly,
the development of In-tensity interferometry discovered by Hanbury
Brownand Twiss (1956a) inspired a new project in
opticalinterferometry. Their basic principle describes
howcorrelations of intensities (not electric fields) can beused to
measure stellar diameters. The importantpoint is that the technique
relies on the correlationbetween the (relatively) low-frequency
intensity fluc-tuations at different detectors, and that it does
notrely on the relative phase of optical waves at thedifferent
detectors. The requirements for the me-chanical and optical
tolerances of an intensity in-terferometer are therefore much less
stringent thanin the case of direct detection schemes as it
isFizeau/Michelson interferometery. First results werereported soon
thereafter (Hanbury Brown and Twiss1956b), leading to the
development of the Narrabriintensity interferometer. With a 188m
longest base-line and blue-sensitivity, this project had a
profoundimpact on the field of optical interferometry, mea-suring
dozens of hot-star diameters (e.g. HanburyBrown et al. 1967a,b,
1970, 1974a,b, Davis et al.1970). The small bandwidths attainable
with in-tensity interferometry limited the technique to
thebrightest stars, and pushed the development of so-called direct
detection schemes, where the light iscombined before detection to
allow large observingbandwidths.
With the advent of lasers in visible and in-frared, the benefits
of radio interferometry havebeen pursued in optical long-baseline
interferometrythrough heterodyne detection (e.g. Gay and
Journet1973, Assus et al. 1979). Radio interferometry func-tions in
a fundamentally different way from opticalinterferometry. Radio
telescope arrays are hetero-dyne, meaning that incoming radiation
is interferedwith a local oscillator signal before detection.
Thesignal can then be amplified and correlated with sig-nals from
other telescopes to extract visibility mea-surements. Optical
interferometers are traditionallyhomodyne, meaning that incoming
radiation is in-terfered only with light from other telescope.
Thisrequires transport of the light to a central station,without
the benefit of being able to amplify the sig-nal. The other
important benefit of radio interfer-ometry is that the required
accuracy for beam re-combination is much more easy to achieve in
radiothen in optical domain. One heterodyne optical in-terferometer
(ISI, see Section 3.1) has been built tooperate at 10m wavelengths.
As for Intensity in-terferometry, the technique is feasible but
with smallbandwidths attainable and limited to bright sources,while
the homodyne optical interferometry allowslarge bandwidths to be
used since the interfered lightis detected directly. Two important
steps towardsmodern optical interferometry have been done at
theObservatoire de Calern, France:
1. Labeyrie (1970) proposed speckle inter-ferometry, a process
that deciphers the diffraction-limited Fourier spectrum and image
features of stel-lar objects by taking a large number of
very-short-exposure images of the same field. This techniquecan be
applied to reconstruct a single image that is
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OPTICAL INTERFEROMETRY I
free of turbulent areas, in essence an image of theobject as it
might appear free of atmospheric turbu-lence.
2. Long-baseline interferometry on a star bydirectly combining
the electric fields before photondetection (direct detection) was
first accomplishedin 1974 by Labeyrie (1975) by using a 12m
baselineat the I2T interferometer.
The following section deals with the funda-mental observables in
astronomical optical interfer-ometry. In Section 3. three basic
designs to achievecoherent combination of light are explained in
moredetails, while in Section 4. the basic
interferometrictechniques are presented. Further development inthe
field is described in Sections 5. and 6. while thepossibilities of
new ground-based and space projectsare discussed in Section 7.
Finally, I summarize themost important conclusions in Section
8.
2. INTERFEROMETRIC OBSERVABLES
2.1. Visibility modulus
The fringe contrast observed with an interfer-ometer is
historically called the visibility and, for thesimple (two-slit)
interferometer can be written as:
V =Imax IminImax + Imin
=Fringe amplitudeAverage intensity
(1)
where Imax and Imin denote the maximum and min-imum intensity of
the fringes.
Generally, the primary observable of an inter-ferometer is the
complex visibility function (the am-plitude and phase), which is
the Fourier transformof the object brigthness distribution (van
Cittert-Zernike Theorem). Let I(u, v, ) denote the spa-tial Fourier
transform of sky brightness I(s, p, ) atthe wavelength , where the
Fourier frequency pairu, v (spatial frequency) is the conjugate of
the spa-tial Cartesian coordinates s, p, respectively. Defin-ing
the complex visibility V (u, v, ) = I(u, v, )/Nwhere N is
normalization factor defined so that|V (0, 0, )| = 1, it
follows:
V = |V (u, v, )|.
The visibility V as defined by Eq. (1) is exactlyproportional to
the amplitude modulus of the imageFourier component corresponding
to the spatial fre-quency (u, v) = ~b/ where ~b is the baseline
vector ~Bprojected onto the plane of the sky.
The phase of the fringe pattern is equal tothe Fourier phase of
the same spatial frequency com-ponent but it cannot be observed
since an incomingplane wave from a source is corrupted as it
propa-gates through the turbulent atmosphere which pro-duces
atmospheric phase delays such that the phaseinformation on the
source is completely lost. Prac-tically only the modulus of the
complex visibilityV = |V (u, v, )| (or squared visibility amplitude
V 2)is an observable quantity.
The corruption of this phase informationhas serious
consequences, since imaging of non-centrosymmetric objects rely on
the Fourier phaseinformation encoded in this intrinsic phase of
inter-ferometer fringes. Without this information, imag-ing cannot
be done except for simple objects such asdiscs or round stars. A
number of strategies havebeen developed to circumvent these
difficulties.
2.2. Relative phase
The technique called phase referencing relieson the fact that
the relative phase can be obtainedif atmospheric and instrumental
noise can be con-trolled or compensated, for example through
ref-erencing by measuring the instantaneous phases offringes using
a reference star (Shao and Colavita1992) or by using the same
source at another wave-length. The technique, called Differential
SpeckleInterferometry (DSI) measures the relative shift ofan object
in different spectral bands, and its appli-cation is based on the
assumption that the shift be-tween two speckle images
(instantaneous distortedimages, see Section 4.1) can be related to
spatialstructures of the object under study. Althoughthe phase of
the spatial coherence function is cor-rupted by instrumental and
atmospheric noise, dif-ferent methods of data processing can be
used toovercome this difficulty. If the two sets of
speckleinterferograms of the same object are recorded
si-multaneously at two different wavelengths, but closeenough for
the Point Spread Function to be the samefor both channels, then
their ensemble average cross-spectrum provides the information
about the rela-tive shift as a fraction of the object size. This
tech-nique was described by Beckers (1982), Jankov et
al.(2001).
2.3. Closure phases
The closure phase measurements tell us aboutthe symmetry of an
object and is only available froma closed triangle of elements,
such as three tele-scopes. Normally, atmospheric fluctuations
disturbthe fringe phase between any two interferometer el-ements,
but these fluctuations cancel out in a closedtriangle.
Considering a 3-telescope array where thephases ij, jk, ki are
intrinsic phases, ij,
jk,
ki
are phases measured on the baselines between tele-scopes (i and
j), (j and k), (k and i) respectively, i,j, k are the atmospheric
errors introduced abovetelescopes (i,j,k) and ij, jk, ki is the
noise, it fol-lows:
ij = ij + i j + ijjk = jk + j k + jkki = ki + k i + ki
Consequently the closure phase:
cl = ij+jk+
ki = ij+jk+ki+ij+jk+ki
is free of atmospheric phase noise (Jennison 1958).
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S. JANKOV
2.4. Bispectrum
When analysing the speckle interferometrydata, one can think of
many virtual sub-apertures(with sizes equal to the coherence length
2 ) spreadacross the full telescope, with fringes forming be-tween
all the sub-aperture pairs. The speckle in-terferometry data can be
reduced using the bispec-trum, which permits a direct inversion
from the es-timated Fourier amplitudes and phases. The
bispec-trum:
Bijk = VijVjkVkiis formed through triple products of the complex
vis-ibilities around a closed triangle, where i, j, k specifiesthe
three aperture locations on the pupil of the tele-scope.
It was discovered that the Fourier phasescould also be estimated
from such data (e.g. Knoxand Thompson 1974), and that the
bispectrumphase is identical to the closure phase. Interest-ingly,
the use of the bispectrum for reconstruct-ing diffraction-limited
images was developed inde-pendently (Weigelt 1977) of the closure
phase tech-niques, and the connection between the approacheswas
realized only later (Roddier 1986, Cornwell1987).
3. INSTRUMENTAL DESIGNS
3.1. Heterodyne interferometer
The principle of a Heterodyne interferometeris to change the
frequency of light at each telescope,carry to the common focus an
intermediate frequencysignal, and combine all these signals in
order to pro-duce fringes. Associated to each telescope is a
laser,and the two beams (laser and star) are then com-bined with a
beam splitter onto an infrared-sensitivephotodiode, which produces
an electric current pro-portional to the light intensity on it. Two
lightwaves, those from the laser and from the star, arevery close
in frequency and while combining twobeams, they beat together to
produce a signal atlower frequency such that it becomes a radio
wave.
To produce interference fringes two radio sig-nals are
overlapped instead of light beams. All ofthe information that is
carried in the visible or in-frared light wave from the star is
still in the newradio signal, the only thing that is different is
thewavelength. The principle of heterodyne interferom-etry is
depicted in Fig. 1.
This method of converting the frequency ofa signal from the star
by mixing it with a fixed-frequency local oscillator is called
heterodyne de-tection. (The term local oscillator refers to
thelaser in this case). Usually, the local oscillatorsare CO2
lasers operating at frequencies 10mwavelengths. The heterodyne
detection suffers frombandwidth limitations as well as additional
noisecontribution from laser shot-noise, which becomesprogressively
worse at higher frequencies. However,one of the main advantages is
that the signals areconverted to radio frequencies and then
combinedelectronically using techniques from radio astronomy.
Fig. 1. Heterodyne interferometer arrangement.The light from an
infrared laser together with infraredlight from a star are
converted to radio frequenciesthrough heterodyne circuits and then
combined elec-tronically. A light wave from a star hits first
tele-scope, and has to travel an extra distance (d) beforeit
arrives at second telescope. In order for the in-terference method
to work properly, the signal fromone telescope should be delayed in
time () until thewavefront reaches the other telescope; only then
twosignals can interfere together.
3.2. Intensity interferometer
The Intensity interferometer is also called thepost-detection
correlation Interferometer since thesignals are correlated after
detection. The inten-sity interferometry technique does not rely on
theactual interference of light rays, thus no fringes areformed.
Instead, the interferometric signal, the de-gree of mutual
coherence, is characterized by the de-gree of correlation of light
intensity fluctuations ob-served at two different detectors. In
practice, the de-gree of mutual coherence is measured using fast
tem-poral correlations between narrow optical wavebandintensity
fluctuations observed by two (or more) tele-scopes separated by a
baseline distance (Fig. 2).
Fig. 2. Intensity interferometer arrangement. Theinterferometric
signal, the degree of mutual coher-ence, is characterized by the
degree of correlation oflight intensity fluctuations observed at
two differentdetectors.
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OPTICAL INTERFEROMETRY I
The principal component of the intensity fluc-tuation is the
classical shot noise which will notdemonstrate any correlation
between the two sepa-rated telescopes. The intensity
interferometric signalis related to a smaller noise component: the
wavenoise. The wave noise can be understood as thebeat frequency in
optical intensity between the dif-ferent Fourier components of the
light reaching thetelescopes. This wave noise will show correlation
be-tween the two detectors, provided there is some de-gree of
mutual coherence between the light receivedat the two
telescopes.
The technique has the benefit of being muchless sensitive to
atmospheric phase fluctuations be-cause the signals input to the
correlator have followedroughly the same path through the
atmosphere. De-lay tracking and local oscillator stability do not
haveto be very accurate. Unlike the direct
(Amplitudeinterferometry) the Intensity interferometry is
wellmatched to blue/visible operation, and scales to longbaselines
with ease (the atmospheric stability re-quirements are a factor of
a million less stringent).
The major drawback of Intensity Interferom-etry is that copious
quantities of light are necessaryto observe the wave noise signal
in the presence ofthe much larger shot noise. Even for brightest
stars,very large light collectors ( 100m2 area) are nec-essary to
have a sufficient statistical strength to ob-serve non-Poisson
fluctuations around the mean pho-ton intensity. Mainly for this
reason, the experimentconducted at Narrabri from 1963 to 1971 was
theonly implementation of a stellar intensity interfer-ometer.
Since then, all stellar interferometers havebeen based on
Fizeau/Michelsons techniques.
3.3. Direct interferometer
In this approach the light from each apertureis carried to a
common focus, combined coherently inorder to detect interferometric
signal (fringes). Thecombination of two (or more) light beams can
bedone in the image plane (Fizeau mode, Fig. 3a)or in the pupil
plane (Michelson mode, Fig. 3b).In the Fizeau mode, the ratio of
aperture diame-ter/separation is constant from light collection to
re-combination in the image-plane (homothetic pupil).In the
Michelson mode, this ratio is not constantsince the collimated
beams have the same diameterfrom the output of the telescope to the
recombinationlens. This means the object-image relationship canno
longer be described as a convolution. However,it is possible to use
a relay optics after the beam-combiner so as to stack two output
pupils on the topof each other with a modulation depending on
theoutput Airy discs distance (Fig. 3c). The main ad-vantage over
the Michelson mode recombinations isconservation of the
object-image convolution relation(Vakili et al. 2004).
Interferometers in which the output pupil isnot a homothetic
image of the input pupil, but thepattern of the subaperture centers
is conserved areknown as densified pupil interferometers
(Labeyrie1996). In such an arrangement, the interferometer
forms a true image of small sources at the combinedfocus. This
is particularly attractive for arrays con-sisting of many
telescopes, because the desired infor-mation can be obtained after
deconvolution with theknown point spread function.
Fig. 3. Comparison of three different beam-combinations for an
optical stellar interferometer.(A) and (B) the classical Fizeau
versus Michelsonbeam-combinations, (C) instead of superimposing
theAiry patterns from the telescopes, it is possible to usea relay
lens after the beam-combiner so as to stackthe two output pupils on
the top of each other with amodulation depending on the output Airy
discs dis-tance.
The disadvantage of the Michelson mode is avery narrow field of
view compared to the Fizeaumode, but for diluted optical arrays it
is extremelydifficult to construct an interferometer with
imageplane beam recombination. For this reason Fizeaumode is more
typically implemented on a common-mount interferometers. Typical
example of such aninterferometer designed to create high resolution
im-ages over a wide field of view is the Large Binocu-lar Telescope
(see Section 6.7), a single-mount struc-ture resulting in a fixed
entrance pupil for which theFizeau mode recombination is
adapted.
The Michelson mode interferometry is usedin long-baseline
interferometers made of independenttelescopes. Beam transfer optics
and delay lines al-low it to maintain equal path length between
thetelescopes as the object is tracked across the sky.
The real art of developing interferometers isto combine the
beams in phase with each other afterthey have traversed exactly the
same optical pathfrom the source through each of the telescope
downto the beam combination point. The paths are madeequal by
adjusting the position of mirrors in the opti-cal delay-line (Fig.
4.) that corrects the drift inducedby the diurnal rotation of the
tracked star. Due tothe long path lengths between the telescope and
cen-tral beam combining facility, a significant
differentialchromatic dispersion occurs if the light is
propagat-ing in the air. In order to combine broad bandwidths,one
must either transport the light through a vac-uum or construct a
dispersion compensator (Tango1990).
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S. JANKOV
Fig. 4. Michelsons interferometer. The light fromeach aperture
is carried to a common focus, combinedcoherently in order to detect
interferometric signal(fringes). Beam transfer optics and delay
lines al-low it to maintain equal path length between the
tele-scopes as the object is tracked across the sky.
In addition, mechanical constraints on the in-strument, errors
on the pointing model, thermaldrifts, various vibrations and
atmospheric turbulencemake the null Optical Path Difference (OPD)
pointchanging. The error on the OPD must be less thanthe coherence
length defined by: lc =
2
where is the mean wavelength observed and is thespectral
interval. Fringes are searched by adjustingthe delay-line position,
and this real-time control iscalled fringetracking.
After being collected by the telescopes andproperly delayed, the
light must be directed to acentral facility for beam combination.
The interfer-ogram can then be recorded, as long as the
scanningtakes place faster than the atmospheric coherencetime c =
c/.
4. INTERFEROMETRIC TECHNIQUES
4.1. Speckle Interferometry
Speckle interferometry is a method permit-ting the extraction of
spatial information from two-dimensional images at scales down to
the diffractionlimit of the telescope, in spite of severe blurring
in-troduced by atmospheric turbulence. The image of astar obtained
through a large telescope looks speck-led or grainy because
different parts of the image areblurred by small areas of
turbulence in the earthsatmosphere. The image obtained for an
unresolvedpoint source depends on the exposure time. For
longexposures, the image is blurred to a seeing disk 1arcsec. For
exposures shorter than the coherencetime (tc 10ms for the
atmosphere in the optical;100ms for the infrared), a group of
bright speckles isobtained approximately of the size of the Airy
disk.In that way the speckle interferometry freezes theatmosphere,
and gets an instantaneously distortedimage as presented in Fig.
5.
Fig. 5. Speckle, an instantaneously distorted stellarimage.
In 1970, the French astronomer AntoineLabeyrie showed that
information could be ob-tained on the high-resolution structure of
the objectfrom the speckle patterns using the Fourier
analysis(speckle interferometry). In the 1980s, the meth-ods which
allowed images to be reconstructed inter-ferometrically from these
speckle patterns were de-veloped using a sequence of short-exposure
snapshotsto obtain images at a telescopes diffraction limit.
With existing large telescopes, speckle tech-niques thus permit
resolution at spatial scales of0.025 arcseconds rather than the 1
to 2 arcsecondsassociated with classical techniques. These meth-ods
are also characterized by enhanced measure-ment accuracy of the
separation of closely spacedobjects seen through the turbulent
atmosphere. Thespeckle interferometry system incorporates an
inten-sified charge coupled device array as the primaryimaging
detector and an autocorrelator as a highspeed data reduction
processor operating at videorates. Also, the reduction of the
speckle patterns canbe done a posteriori involving computer
processing.
4.2. Aperture Masking Interferometry
Aperture Masking Interferometry is a form ofspeckle
interferometry, allowing to reach the maxi-mum possible resolution
at ground-based telescopeswith large diameters to produce
diffraction limitedimaging.
In the aperture masking technique, the bis-pectral analysis
(speckle masking) method is typi-cally applied to data taken
through masked aper-tures, where most of the aperture is blocked
off andlight can only pass through a series of small
holes(subapertures) so that the array of holes acts as aminiature
astronomical interferometer. The aperturemask removes atmospheric
noise from these measure-ments, allowing the bispectrum to be
measured moreprecisely than for an un-masked aperture. For
sim-plicity the aperture masks are usually either placedin front of
the secondary mirror (e.g. Tuthill et al.2000, see Fig. 6) or
placed in a re-imaged apertureplane (e.g. Young et al. 2000).
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OPTICAL INTERFEROMETRY I
Fig. 6. Desired sparse pupil (left) and the aperturemask used to
generate it (right). The required pupilis shown as the set of black
spots and is superimposedon a scaled version of the segmented 10m
Keck pri-mary mirror (hexagonal segments). From Tuthil etal.
(2000).
The masks are usually categorised eitheras non-redundant or
partially redundant. Non-redundant masks consist of arrays of small
holeswhere two pairs of holes have not the same sepa-ration vector
(the same baseline). Each pair of holesprovides a set of fringes at
a unique spatial frequencyin the image plane.
Partially redundant masks are usually de-signed to provide a
compromise between minimis-ing the redundancy of spacings and
maximising boththe throughput and the range of spatial frequen-cies
investigated. Although the signal-to-noise ofspeckle masking
observations at high light level canbe improved with aperture
masks, the faintest limit-ing magnitude cannot be significantly
improved forphoton-noise limited detectors so that the
principallimitation of the technique is that it is limited
torelatively bright astronomical objects.
4.3. Aperture Synthesis
The direct closure phase measurements to-gether with the
measurements of visibility amplitudeallow one to reconstruct an
image of any object us-ing three or more independent telescopes.
This tech-nique has been successfully demonstrated by Bald-win et
al. (1996) in the visible band at the Cam-bridge optical
aperture-synthesis telescope (COAST,see Section 5.3). Each pair of
telescopes in an arrayyields a measure of the amplitude of the
spatial co-herence function of the object at a spatial frequency~b/
which corresponds to one point in (u, v) plane.
In order to make an image from an interferom-eter, one needs
estimates of the complex visibilitiesover a large portion of the
(u, v) plane, both the am-plitudes and phases. The (u, v) points
correspondingto a snapshot projection of the baseline are
sampledfrom each available baseline. It is also useful to use
atechnique called super-synthesis: the (u, v) planeis swept during
an observation lasting several hours,due to Earth rotation. After a
large variation of hourangle, several visibility moduli are,
therefore, measu-
Fig. 7. uv-plane coverage for the observations ofBaldwin et al.
(1996). Each point corresponds to anindividual visibility
measurement.
red at different positions in (u, v) plane as presentedin Fig.
7.
Reconstruction of complex images involves theknowledge of
complex visibilities but the phase of avisibility may be deduced
from closure-phase thathas been applied in optical interferometry.
The pro-cess after data acquisition consists of phase calibra-tion
and visibility phase reconstruction from closure-phase terms by a
technique similar to bispectrumprocessing, so that at least three
configurations ofthree telescope interferometer is needed.
Interferom-eters with two apertures have limited possibilities
forimage reconstructions due to absence of the phasevisibility
recovering but in that case different spec-tral channels can be
employed for differential visibil-ity measurements (between the
continuum and spec-tral line) and use of data from one part of the
spec-trum such as the continuum to calibrate another part(Jankov et
al. 2001, 2002, 2003).
4.4. Nulling Interferometry
This technique, the nulling interferometry,holds a great promise
for the detection and char-acterization of Earth-like extrasolar
planets at mid-infrared wavelengths, because the light from
theparent star arriving on axis is completely rejected(Bracewell
1978). The basic premise of nulling in-terferometry is conceptually
quite simple: combinethe starlight that arrives at a pair of
telescopes sothat at the centre of the image the two waves
areexactly pi out of phase producing a dark centralfringe and
effectively making the central star disap-pear. A deep destructive
fringe is to be placed acrossthe star, but light coming from
sources that are off-set slightly from it (i.e. from planets) gets
addedrather that subtracted so that even though the staris
completely blanked out, planets at the right loca-tions (near a
constructive fringe) are not attenuated(Fig. 8).
An alternative to the nulling Michelson inter-ferometer is the
use of a transparent phase-shiftingmask in conjunction with a
densified pupil interfer-ometer (Boccaletti et al. 2000). The basic
idea of
7
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S. JANKOV
Fig. 8. Nulling Interferometry Principle. Thestarlight that
arrives at a pair of telescopes so that adeep destructive fringe is
placed across the star (mak-ing it disappear), but the light coming
from a planetgets added rather that subtracted.
this concept is to introduce a pi phase shift over thecentral
part of the point spread function produced bythe densified pupil; a
careful adjustment of the sizeof the phase mask together with pupil
apodizationin a coronographic setup can in principle provide
anexcellent nulling performance.
Nulling interometry can be performed fromthe Earth, especially
when imaging Jupiter sizedplanets, but the best place for viewing
earth-sizedplanets close to their suns is from a space-based
plat-form. Both NASA and ESA have plans to launchinfrared nulling
interferometry telescopes into spaceover the next decades (see
Section 7).
5. PAST INTERFEROMETERS
In this section, only the most important inter-ferometers that
produced significant results in thepast are presented:
5.1. I2T
The I2T (Interferome`tre a` 2 Telescopes), Ob-servatoire Cote
DAzur, Plateau de Calern, France(Fig. 9) is the interferometer
where the first directinterference fringes between separate
telescopes wereobserved and reported by Labeyrie (1975).
Fig. 9. The I2T (Interferome`tre a` 2 Telescopes),Plateau de
Calern, France.
I2T made observations in the visible (e.g.,Blazit et al. 1977,
Thom et al. 1986) and pioneeredinterferometry at near-IR
wavelengths (di Benedettoand Conti 1983, di Benedetto 1985).
5.2. MARK I, II and III
The Massachusetts Institute of Technologyand the Naval Research
Laboratory built and op-erated a series of prototype
interferometers, namedthe Mark I, Mark II, and Mark III located on
Mt.Wilson, California, US.
Active fringe tracking was first demonstratedin 1979 by the Mark
I interferometer (Shao andStaelin 1980), while the most popular
architecturetoday for the moving delay line is based on the
so-lution implemented by the Mark III interferometer(Shao et al.
1988).
The Mark III was specifically designed to per-form wide-angle
astrometry, but a variable baselinethat could be configured from 3m
to 31.5m providedthe flexibility needed for a variety of
astronomicalprograms. Because a full computer control of
thesiderostats and delay lines allowed automated acqui-sition of
stars and data acquisition, it could observeup to 200 stars in a
single night, until 1993, whenthe interferometer stopped the
operation.
5.3. COAST
The COAST (Cambridge Optical ApertureSynthesis Telescope)
operated from 1991 until 2005at Cambridge University, England. It
was plannedas a coherent array of four telescopes operating inthe
red and near infra-red, using Michelson interfer-ometry on
baselines of up to 100m to give imageswith a resolution down to
1mas.
The number of telescopes (40cm apertures)has been brought up to
five, and light from fourof them could be interfered
simultaneously. Switch-ing between two four-telescope arrays
allowed takingdata on nine baselines and seven closure
trianglesduring a single night. Observations could be carriedout in
the visible (R,I) or near-infrared (J,H bands).It was the first
instrument of its kind to exploit thetechniques of aperture
synthesis and closure phaseat optical or infra-red wavelengths,
producing thefirst images from an optical aperture synthesis
array(Baldwin et al. 1996).
5.4. IOTA
The IOTA (Infrared Optical Telescope Ar-ray) operated from 1993
to 2006 and was an in-terferometer of Smithsonian Astrophysical
Observa-tory on Mt. Hopkins, Arizona, US (Fig. 10), withthree 45cm
telescopes (enabling closure phase obser-vations) and baselines
from 5m up to 38m. Observa-tions have been carried out in the
visible or near-IR(J,H,K bands). Visibilities with excellent
calibrationcould be obtained with the single-mode fiber systemFLUOR
(Fiber Linked Unit for Optical Recombina-tion, Coude Du Foresto et
al. 1997) which accountsfor a large fraction of the astronomical
results ob-tained with IOTA.
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OPTICAL INTERFEROMETRY I
Fig. 10. The IOTA (Infrared Optical TelescopeArray) on Mt.
Hopkins, US.
5.5. GI2T
The GI2T (Grand Interferome`tre a` 2Telescopes), Observatoire
Cote DAzur, Plateau deCalern, France (Mourard et al. 1994, see Fig.
11)was a successor of the I2T since 1985 and used twoboule
telescopes with 1.5m apertures on a north-south baseline that could
be reconfigured from 10mto 65m.
Fig. 11. GI2T (Grand Interferome`tre a` 2 Telesco-pes), Plateau
de Calern, France.
The GI2T had the capability of performingspectrally resolved
interferometry. After cloture in2006, its beam combination table,
including a ver-satile visible spectrograph and an IR focus has
beenmoved to CHARA array (Mourard et al. 2009).
5.6. MIRA
The MIRA (Mitake Infrared Array), NationalAstronomical
Observatory, Japan, was an ambitiousplan to build a series of
interferometers with in-creasing capabilities. The first phase of
this project(MIRA-I), which consisted of two 25cm telescopeswith
coude optics on a 4m baseline in Tokyo has suc-cessfully been
completed in 1998, and has acquired
stellar fringes. The next step was to built an instru-ment with
a slightly larger siderostats and a 30mbaseline (MIRA-I.2). The
interferometer ceased op-eration in 2007.
5.7. PTI
The PTI (Palomar Testbed Interferometer),NASA JPL, Mt Palomar,
CA, US (Colavita et al.1999, see Fig. 12) was built in 1995 and it
consistedof three 40cm siderostats, up to 110m baselines
whichprovided 3mas resolution in the near-infrared (H andK bands).
It was developed primarily to demonstratethe utility of
ground-based differential astrometry inthe search for planets
around nearby stars, and todevelop key technologies for the Keck
Interferometerand space-based missions.
Fig. 12. PTI aerial figure. PTI is located atopPalomar Mountain,
US, next to the large white domeof the historic 5m Hale
Telescope.
PTI was notable for being equipped with adual-star tracking
system, the first of its kind,which simultaneously tracked
interference fringesfrom a target star and a reference star against
whichthe target was measured. This allowed to cancelsome of the
atmospheric effects of astronomical see-ing and to make very high
precision measurementspossible.
Aside from its role for the technical develop-ment of high
precision dual-star astrometry, thePTI was used mainly for stellar
diameter measure-ments, stellar masses and binary star work. The
in-strument concluded operations in 2009.
6. EXISTING INTERFEROMETERS
6.1. NPOI
The NPOI (Navy Prototype Optical Interfer-ometer), located on
Anderson Mesa near Flagstaff,Arizona, US (see Fig. 13) operates
since 1995 andcombines an imaging array and an astrometric
fa-cility (Armstrong et al. 1998). It can observe in
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S. JANKOV
the visible with 32 spectral channels covering thewavelength
range from 450nm to 850nm. The imag-ing subarray consists of six
movable siderostats withbaseline lengths from 2.0m to 437m. The
array ge-ometry has been optimized for baseline bootstrap-ping to
facilitate the imaging of stellar surface struc-ture. Like the Mark
III, the NPOI uses vacuum delaylines for pathlength compensation.
The four-elementastrometric subarray of the NPOI includes an
exten-sive site metrology system that monitors the motionsof the
siderostats with respect to one another in orderto perform
wide-angle astrometry with more than2mas precision.
Fig. 13. The NPOI (Navy Prototype Optical Inter-ferometer),
located on Anderson Mesa, US.
6.2. SUSI
The SUSI (Sydney University Stellar Inter-ferometer) is an array
of 13 telescopes with 14cmapertures, operating since 1991 at Sydney
UniversityNarrabri, Australia (Davis et al. 1999, see Fig. 14).The
SUSI has two beam-combining systems: Theoriginal blue system was
designed to operate inthe wavelength range 400-540 nm and employs
nar-row bandwidths (typically a few nanometers). Thenewer red
system is designed to work in the range500-950 nm. SUSI makes
observations on a singlebaseline selected from a set of fixed
north-south base-lines with lengths ranging from 5m to 640m and
itcan achieve angular resolutions from 20mas to 50as.
Fig. 14. The SUSI (Sydney University StellarInterferometer),
Narrabri, Australia.
The 640m baseline length, the longest of allinstruments
currently operational or under construc-tion, has been chosen to
resolve a sample of O stars ata wavelength of 450 nm. The focus of
SUSIs obser-vations is on improving our understanding of
stellarastrophysics including: single stars (measuring ef-fective
temperatures, radii and luminosities), binarystars, (as for single
stars, plus measuring distancesand masses), variable stars (e.g.
Cepheids and Mi-ras), and emission line stars (e.g. Be and
Wolf-Rayetstars).
6.3. ISI
The ISI (Infrared Spatial Interferometer),University of
California at Berkeley, US is locatedclose to the CHARA array and
the former site ofthe Mark III on Mt. Wilson. It started operation
in1990 as two-telescope interferometer but presentlyconsists of
three 1.65m telescopes observing in themid-infrared. The telescopes
are fully mobile andtheir current site on Mount Wilson allows for
place-ments as far as 70m apart providing a resolution of3mas at
11m. On July 2003, the ISI recorded itsfirst closure phase aperture
synthesis measurements.
The Fig. 15 shows three ISIs telescopes all ina line which were
used for initial testing purposes,with 4m, 8m, and 12m baselines.
However, they canbe moved in such a way that they form a
triangleand three baselines at three different angles can
bemeasured simultaneously providing the closure
phasemeasurments.
Fig. 15. The ISI (Infrared Spatial Interferometer),Mt. Wilson,
US.
The interferometer operates at wavelengthsbetween 9m and 12m and
the stellar radiation ismixed with the output of a CO2 laser, which
acts asthe local oscillator. Observations have been carriedout with
baseline lengths up to 56m. One of manyadvantages of the ISIs
narrow heterodyne detectionbandwidth is that one can tune the
detection wave-length to be deliberately in or out of known
spectrallines, allowing for interferometry on spectral lines tobe
carried out.
Interest of the ISI science has mainly been fo-cused on the
study of evolved stars and dust shells.However, an interferometer
like the ISI is well suited
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OPTICAL INTERFEROMETRY I
for making very precise measurements of positionsof stars and
particularly to measure positions of in-frared stars which are
hardly detectable in visiblelight. Investigations in the field of
astrometry shouldhelp to tie the astronomical reference frames
togetherwith that established in the radio, infrared and
visi-ble.
6.4. CHARA
The CHARA array, named after Center forHigh Angular Resolution
Astronomy, Georgia StateUniversity is located on Mt Wilson, US (see
Fig.16). It consists of 6 telescopes of 1m in diameterarranged in a
Y-shaped configuration with baselinesranging from 30m to 330m.
First fringes on a sin-gle baseline have been obtained in November
1999and commissioning of the full array continues. Lightfrom the
individual telescopes is conveyed throughvacuum tubes to the
central Beam Synthesis Facil-ity in which the six beams can be
combined together.When the paths of the individual beams are
matchedto an accuracy of less than one micron, after the
lighttraverses distances of hundreds of meters, the Arraythen acts
like a single coherent telescope for the pur-poses of achieving an
exceptionally high angular res-olution. The Array is capable of
resolving details assmall as 200as.
Fig. 16. The facilities of the CHARA (Centerfor High Angular
Resolution Astronomy) Array to-gether with the existing facilities
of historic MountWilson Observatory. The Beam Synthesis Facility,an
L-shaped structure of the length of a football field,can be seen
close to the 2.5m telescope dome.
CHARA has entered into several collabora-tions with groups
offering unique instruments ortechnologies for enhanced
performance. These in-ternational collaborations have brought a
significantadded value to the science capabilities of the
CHARAArray. At present these collaborations include ajoint
observing collaboration with the Observatoirede Paris through the
FLUOR instrument which hasbeen moved from the IOTA interferometer
(Section5.4) and upgraded for the CHARA Beam Synthe-sis Facility. A
collaboration with the University of
Michigan has led to the development of an imagingbeam combiner
(MIRC) which has already producedthe first images of stellar
surfaces and close binarystars. A joint project with the University
of Syd-ney has led to the development of a second beamcombiner with
significantly improved sensitivity tofainter objects while also
providing measurements ofvery high precision. Finally, an agreement
with Ob-servatoire de la Cote dAzur has brought the thirdnew beam
combiner (VEGA) to the Array capableof providing spectroscopic and
polarimetric channelsfor high resolution work.
The first four-telescope fringes with VEGABeam Combiner at the
CHARA Array were obtainedwith CHARA (MIRC used as the infrared
fringetracker) in October 2010.
6.5. KI
The KI (Keck Interferometer), NASA JPL, onMauna Kea, Hawai, US
(Fig. 17) which consists oftwo 10m Keck telescopes obtained first
fringes withfull-apertures on March 2001.
Fig. 17. Two 10m telescopes of the KI (Keck In-terferometer) on
Mauna Kea, Hawai, US.
The 10m telescopes are equipped with high-order adaptive optics
systems, which provide goodcorrection in the near-IR and excellent
wavefrontquality at 10m. It enables the implementation ofa nulling
beam combiner, which will be used to char-acterize exo-zodiacal
emission around nearby main-sequence stars. The standard science
modes of theinstrument are:
1. V2 mode, high sensitivity visibility ampli-tude measurements
in the near infrared (H and K).
2. Nulling interferometry in N-band (8 -12m) to suppress light
from central star and toreveal faint extended emission around it.
Key sci-entific goals are to characterize exo-zodiacal emis-sion
around Sun-like stars to inform Terestial PlanetFinder (see Section
7.3) mission design.
3. Differential Phase, multi-color fringe phasemeasurements
between 2-5m, with direct HotJupiter detection (orbital parameters
and masses),as the key scientific goal.
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S. JANKOV
Recently, the Keck interferometer was up-graded to do
self-phase-referencing (SPR) assistedK-band spectroscopy at R 2000.
This SPR modeas currently being implemented as the ASTrometricand
phase-Referenced Astronomy (ASTRA) projectwill provide phase
referencing and astrometric ob-servations at the Keck
Interferometer, leading to en-hanced sensitivity, and the ability
to monitor orbitsat an accuracy level of 30-100as, high spectral
res-olution observation of Young Stellar Objects (YSO),sensitive
Active Galactic Nuclei (AGN) observations,and astrometric (known
exo-planetary systems char-acterization and galactic center general
relativity)capabilities. With the first high spectral
resolutionmode now offered to the community, this contribu-tion
focuses on the progress of the dual field andastrometric modes.
It is planned that the KI will be upgradedwith four new 1.8m
outrigger telescopes in order tomake possible imaging mode. A large
fraction of theobserving time available with the outriggers
shouldbe devoted to an astrometric search for extrasolarplanets
while the combination of all six telescopesshould result in a
sensitive imaging array.
6.6. VLTI
The VLTI (Very Large Telescope Interferome-ter), European
Southern Observatory Paranal, Chile,obtained first fringes on the
sky in March 2001, withsiderostats, and presently it allows the
interferomet-ric combination of the four 8.2m telescopes (Fig.
18),augmented by four movable 1.8m auxiliary telescopes(Fig.
19).
The first generation of instruments included acommissioning
instrument based on a two-telescopenear-IR beam combiner VINCI (now
decommis-sioned) as well as actually operating instrumentsMIDI,
AMBER and PRIMA:
MIDI (MID-infrared Interferometric instru-ment) is the
mid-infrared (N-band = 8 to 13m)combiner which combines two beams,
either from theVLT main 8.2m Unit Telescopes (UTs) or from the
Fig. 18. Four 8.2m main Unit Telescopes ofthe VLTI (Very Large
Telescope Interferometer),Paranal, Chile.
Fig. 19. Four 1.8m Auxilary Telescopes ofthe VLTI (Very Large
Telescope Interferometer),Paranal, Chile.
1.8m Auxiliary Telescopes (ATs), to provide visibil-ity moduli
in the (u, v) plane. MIDI features spectro-scopic optics to provide
visibilities at different wave-lengths within the N band and with
spectral resolu-tion R=30 with the prism, R=230 with the grism
(acombination of a diffraction grating and prism). UTadaptive
optics guarantees diffraction-limited imagequality on MIDI, for
targets brighter than V = 17m.These unique capabilities allowed an
access to vari-ous observing programs including observation of
theGalactic Center, analysis of the nuclei of the AGNgalaxies,
particularly stellar and circumstellar ob-jects, as well as solar
system objects.
AMBER is the near-infrared instrument forthe VLTI, which offers
the possibility of combiningtwo or three beams from either the UTs
or the ATs.With spectral resolution up to 10 000, high
visibilityaccuracy and the ability to obtain closure phases,AMBER
offers the means to perform high qualityinterferometric
measurements in the J,H,K bands (1-2.5m range). Angular resolution
is set by the max-imum available baseline, which is about 200
metersfor the ATs and about 130 meters for the UTs. Ac-cordingly,
the limit is about 2mas for the ATs, andabout 3mas for the UTs, in
the K band. These de-sign characteristics, coupled to the VLT
interferome-ter potential, open up the access to investigations
ofseveral classes of objects, from stellar to extragalac-tic
astronomy including, for instance, the solar sys-tem objects,
low-mass stars, Cepheids, microlensing,hot exo-planets, the
galactic center, nearby galax-ies, and Young Stellar Objects,
Luminous Blue Vari-ables, Be stars, Novae, Wolf-Rayet and
Post-AGBstars for which the results have been already
pub-lished.
PRIMA (Phase-Referenced Imaging andMicro-arcsecond Astrometry)
is a system designedto enable simultaneous interferometric
observationsof two objects that are separated by up to 1 ar-cmin,
without requiring a large continuous field ofview. PRIMA improves
the sensitivity by using twoobjects (a bright star guide near a
fainter targetobject) to bring the corrections to the
atmosphericturbulences. Depending on the operation mode,
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OPTICAL INTERFEROMETRY I
PRIMA can be used either to measure the angu-lar separation
between the two objects (astrometrymode) or to produce images of
the fainter of the twoobjects using a phase reference technique
(imagingmode). With the limiting magnitude K 19m withthe UTs and K
15m with the ATs and the angulardistance between the targets that
can be evaluatedwith a resolution of 10as, the objects such as
extra-solar planets, galactic center, Active Galactic Nucleiand
extragalactic objects as well as gravitationalmicro-lensing are the
targets for PRIMA.
In addition to these ESOs instruments thereis a possibility to
bring and install the visiting instru-ments. One of the visiting
instruments PIONIER,developed at LAOG in Grenoble, complements
ex-isting AMBER and MIDI instruments. While AM-BER has previously
combined the light from threeof the telescopes at the VLTI, the
light coming fromthe four 1.8-metre Auxiliary Telescopes has been
suc-cessfully combined for the first time, in October 2010.This is
an important step towards unleashing the fullpotential of the VLTI
to use multiple telescopes to-gether.
The second generation of VLTI instrumentsGRAVITY, VSI and
MATISSE are under develop-ment:
GRAVITY will coherently combine the lightfrom all four UTs. It
will provide near infrared adap-tive optics assisted precision
narrow-angle astrom-etry (10as) and interferometric phase
referencedimaging (4mas) of faint objects (K=20m). It willprecisely
measure the fringe visibility and phase ofthe fringe pattern to
acquire spatial knowledge of theobserved object. With an accuracy
of 10as, GRAV-ITY will be able to study motions to within a
fewtimes the event horizon size of the Galactic Centermassive black
hole and potentially test the GeneralRelativity in its strong field
limit. It will be able todetect intermediate mass black holes
throughout theGalaxy by their gravitational action on
surroundingstars. It will directly determine masses of
exo-planetsand brown dwarfs, and trace the origin of protostel-lar
jets. Also it will be capable of spatially resolvingcoherent gas
motions in the broad line regions of Ac-tive Galactic Nuclei in
external galaxies. Through itshigh performance infrared wavefront
sensing system,GRAVITY will open up deep interferometric imag-ing
studies of stellar and gas components in dusty,obscured regions,
such as obscured active galacticnuclei, dust-embedded star forming
regions, and pro-toplanetary disks.
VSI (V(LTI) Spectro-Imager) was proposedas a second-generation
instrument in order to pro-vide the community with
spectrally-resolved, near-infrared images at angular resolutions
down to1.1mas and spectral resolutions up to R = 12 000.Targets as
faint as K = 13m will be imaged withoutrequiring a brighter nearby
reference object whilefainter targets can be accessed if a suitable
refer-ence is available. The unique combination of
high-dynamic-range imaging at high angular resolutionand high
spectral resolution enables a scientific pro-gram which serves a
broad user community and atthe same time provides the opportunity
for break-throughs in many areas of astrophysics. The highlevel
specifications of the instrument are derived from
a detailed science case based on the capability toobtain
milli-arcsecond resolution images of a widerange of targets
including: probing the initial condi-tions for planet formation in
the AU-scale environ-ments of young stars; imaging convective cells
andother phenomena on the surfaces of stars; mappingthe chemical
and physical environments of evolvedstars, stellar remnants, and
stellar winds, and disen-tangling the central regions of active
galactic nucleiand supermassive black holes. The VSI will
providethese new capabilities using technologies which havebeen
extensively tested in the past. It is quite flexi-ble instrument
and should be able to make maximumuse of new infrastructure as it
becomes available; forexample, by combining up to 8 telescopes,
enablingrapid imaging through the measurement of up to
28visibilities in every wavelength channel within a fewminutes. The
current studies are focused on a 4-telescope version with an
upgrade to a 6-telescopeone.
MATISSE (Multi AperTure mid-InfraredSpectroScopic Experiment) is
foreseen as a mid-infrared spectro-interferometer combining up to
fourUTs or ATs beams. It will measure closure phase re-lations and
thus offer an efficient capability for imagereconstruction. In
addition to the N band, the MA-TISSE will open three new observing
windows at theVLTI: the L, M, and Q bands, all belonging to
themid-infrared domain. Furthermore, the instrumentwill offer the
possibility to perform simultaneous ob-servations in separate bands
while providing severalspectroscopic modes. The unique performance
of theinstrument is related to the existence of the four UTsthat
permits to push the sensitivity limits to valuesrequired by
selected astrophysical programs such asthe study of Active Galactic
Nuclei and protoplane-tary discs. Moreover, the evaluated
performance ofMATISSE is linked to the availability of ATs whichare
relocatable in position in about 30 different sta-tions allowing
the exploration of the Fourier planewith up to 200 meters baseline
length. Key scienceprograms using the ATs cover for example the
forma-tion and evolution of planetary systems, the birth ofmassive
stars as well as the observation of the high-contrast environment
of hot and evolved stars. Insummary, MATISSE can be seen as a
successor ofMIDI by providing imaging capabilities in the en-tire
mid-infrared accessible from the ground. Theextension of MATISSE
down to 2.7m as well as itsgeneralization of the use of closure
phases makes italso a successor of AMBER.
6.7. LBTI
The LBTI (Large Binocular Telescope Inter-ferometer) on Mt.
Graham, Arizona, US (Fig. 20)consists of two 8.4m telescopes
mounted side by sidein a single alt-azimuth mount, with a 14.4m
center-to-center spacing.
This configuration offers some unique capabil-ities for
interferometry, as it lends itself to Fizeaubeam combination with a
wide field of view andlow thermal background. It is a fully
cryogenic in-strument devoted to near and mid-infrared
observa-tions (2-20m). Among the unique characteristics ofthe
telescope are the two adaptive secondary mirrorswhich enable
diffraction limited seeing. The system
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S. JANKOV
Fig. 20. The LBT (Large Binocular Telescope) onMt. Graham,
US.
is designed for high spatial resolution, high dynamicrange
imaging and nulling interferometry. It con-sists of an universal
beam combiner (UBC) and threecamera ports, one of which is
populated currently bythe Nulling and Imaging Camera (NIC).
One of the key projects for the LBTI is to sur-vey nearby stars
for debris disks down to levels whichmay obscure detection of
Earth-like planets. A sur-vey of approximately 80 nearby stars that
is essentialas a ground-based precursor of the space missions
(asfor example TPF/Darwin, see Sections 7.3 and 7.4).A survey will
identify zodiacal disks, which signal theexistence of a planetary
system for follow up studywith TPF/Darwin. On the other hand, the
surveywill also weed out those systems with very intense zo-diacal
emission, which can overwhelm the signal ofterrestrial planets.
Finally, the survey can also ad-dress the question of zodiacal
emission strength as afunction of spectral type and age of the
central star.
In addition, the consortium of German insti-tutes and
observatories led by the Max Planck Insti-tute for Astronomy in
Heidelberg together with theIstituto Nazionale di Astrofisica in
Arcetri, is devel-oping the LINC ((L)BT (IN)terferometric
(C)amera)instrument. The LINC will combine the radiationfrom the
two 8.4m primary mirrors in Fizeau modewhich allow true imagery
over a wide field of view.The beam combiner will operate at
wavelengths be-tween 0.6 and 2.4m. When coupled with the ad-vanced
adaptive optics system of the LBT, the LINCinstrument will allow
the interferometric measure-ments over a field of approximately
10-20 arcsecondssquare. Ultimately, after improvement of the
Adap-tive Optics, which will allow better performance overa wider
field of view, the instrument will be knownas NIRVANA, the
(N)ear-(IR) / (V)isible (A)daptiveI(N)terferometer for (A)stronomy.
The main scien-tific goals of LINC/NIRVANA are YSOs environ-ments,
circumstellar disks and outflows, formationof binary stars, stellar
clusters, the compact H IIregions, astrometry for extra-solar
planets, the so-lar system minor bodies, the galactic center and
thegalaxies at z 1 2.
In October 2010 the LBTI acquired its firstfringes on the sky
and will be interferometricallyfunctional in 2011 once the adaptive
optics and theLBTI are commissioned.
6.8. MROI
The MROI (Magdalena Ridge Observatory In-terferometer) is
presently under construction andshould be operational in years to
come. It is an inter-national scientific collaboration between New
MexicoTech (US) and the University of Cambridge (UK).The
Interferometer Projects mission is to developa ten-element imaging
interferometer to operate atwavelengths between 0.6m and 2.4 m with
base-lines from 7.5m to 340m. Its primarily technical andscientific
goals are to produce model-independent im-ages of faint and complex
astronomical targets atresolutions over 100 times that of the
Hubble SpaceTelescope in order to study: star and planet
for-mation, stellar accretion and mass loss, and activegalactic
nuclei.
The 2.4-meter telescope achieved first light onOctober 2006 and
commenced operations on Septem-ber 2008. The work at the
interferometer site startedin September 2006 and the Beam Combining
Facil-ity was completed in early 2008. Presently, the tele-scopes
are nearing their first factory acceptance tests.When finished, the
interferometer will be upgradedwith ten telescopes, each
approximately 1.4 metersin diameter. The telescopes making up the
interfer-ometer will be spaced by distances up to 400 metersand
will be optically linked in order to make aper-ture synthesis
imaging. This setup will simulate theresolving power of a single
telescope up to 400 me-ters in diameter. As a result of the large
numberof telescopes in the array, the interferometer will beable to
make accurate images of complex astronom-ical objects.
7. FUTURE INTERFEROMETERS
7.1. OHANA
The OHANA (Optical Hawaiian Array forNanoradian Astronomy) is
the project of a hecto-metric fibered array at Mauna Kea, Hawaii.
Theconcentration of large apertures on Mauna Kea (twoKeck, Gemini,
Subaru, CFHT, UKIRT) lends itselfto ideas for interferometric
combination of these sixtelescopes (Mariotti et al. 1996). By
taking advan-tage of the existing adaptive optics systems, such
asalready installed on telescopes, this would provideexcellent
sensitivity and unprecedented imaging ca-pabilities with baselines
ranging from 75m to 800m(the largest baseline of OHANA
(Subaru-Gemini)).This yields resolutions of 0.25mas and 0.5mas at
1and 2 microns respectively, therefore opening thepossibilities for
research in many fields as for VLTI,but with four times better
spatial resolution.
The OHANA project has carried out initialexperiments to couple
the light from Mauna Keatelescopes into a singlemode pair of 300m
fibres andachieved first fringes with small prototype interfer-
14
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OPTICAL INTERFEROMETRY I
ometer on July 2010, the first step in a plan to linkthe giant
telescopes of Hawaii into a powerful opticalinterferometer.
7.2. Interferometry at Antartica
The Dome C at Antartica (altitude 3300m) isconsidered as a
perfect site for astronomical inter-ferometry due to low wind
speeds, little atmosphericturbulence, negligible seismic activity
and very sta-ble climate. For this reason, several
interferometricprojects are under development, most notably theAPI
(Antartic Planet Interferometer, Swain et al.2003) and KEOPS
(Kiloparsec Explorer for OpticalPlanet Search, Vakili et al.
2003).
The API should be built in two phases. Phase1 is proposed to
consist of a two-element active fringetracking interferometer using
two 50cm diametertelescopes operating in a direct visibility (V2)
modewith baselines of up to 400 metres. The H,K,L,Mbands (1.4 to
5.4m) are to be targeted to give max-imum resolution and
sensitivity at up to 104 spectralresolution. In Phase 2, the full
API, will consist ofseveral 2m diameter telescopes and will be
capable ofoperating in a variety of modes. It will be a
partic-ularly powerful tool for the study of extrasolar plan-ets,
protoplanets, Young Stellar Objects, and ActiveGalactic Nuclei
accretion disks.
The KEOPS is proposed as three concentricrings of 1 to 2m
diameter telescopes covering a di-ameter of 650 metres. The 36
co-phased primarytelescopes would have an equivalent collecting
areaof at least 60m2 and a resolution of a milli-arcsecondat 10m.
KEOPS is designed for direct imaging ofexo-planets in the thermal
infrared, as well as forastroseismology studies.
7.3. TPF
The TPF (Terrestrial Planet Finder) Interfer-ometer is a concept
for a formation flying interferom-eter to search for Earth-like
planets around nearbystars, planet studies and signs of life. TPF
was be-ing developed as a possible collaboration with ESAsDarwin
mission but the project was recently can-celled.
The NASAs 2010 Decadal Survey recom-mended a continuation of
technology developmentfor nulling interferometry under the title of
NewWorlds Technology Development Program and now,although TPF no
longer exists as a planned mis-sion within NASA, some technology
work continues.Efforts in 2009 and 2010 have included
broadbandnulling with the Adaptive Nuller and demonstrationswith
the Planet Detection Testbed. One of the pos-sible configurations
of a space interferometer (an op-tion considered for Terrestial
Planet Finder) is pre-sented in Fig. 21.
At present, NASA continues to work with Uni-versity scientists
and industry to develop the TPFtechnology. Meanwhile, the Kepler
mission will col-lect information on the number of possible
nearbyEarthlike planets, while the Keck Interferometer andLarge
Binocular Telescope Interferometer projectswill measure, for this
purpose, the dust surroundingthose stars.
Fig. 21. A possible configuration of a space
inter-ferometer.
7.4. Darwin
The Darwin mission was planned as a con-stellation of four or
five spacecrafts, with primarygoal to look for Earth-like planets
and analyse theiratmospheres for chemical signatures of life.
Onespacecraft would have acted as a central communica-tions
station. The other three/four would have func-tioned as light
collectors, redirecting light beams tothe central station. The
constellation was also in-tended to carry out high-resolution
imaging usingaperture synthesis, to provide images of celestial
ob-jects with unprecedented detail. Instead of orbit-ing Earth,
Darwin was proposed to be placed be-yond the Moon. At a distance of
1.5 million kmfrom Earth, in a direction opposite to that of
theSun, Darwin would have operated from the secondLagrange point
L2. The Darwin mission was pro-posed in 2007 as a concept for Phase
A Study withinESAs Cosmic Vision, which includes missions be-yond
2015. It was not chosen for further study, but asimilar project
should be re-proposed to ESA. What-ever the final design of the
telescopes, it is clear thatthe TPF/Darwin technology will be a
pathfinder forfuture micro and nano-arcsecond resolution
instru-ments at infrared and other wavelengths.
8. CONCLUSION
After decades of development, optical interfer-ometry is now
beginning to play a major role in astro-physics. The emergence of
large interferometers, inparticular the Very Large Telescope
interferometer,CHARA, the Keck Interferometer and Large Binocu-lar
Telescope Interferometer promise to revolutionizethe impact of high
resolution observations in manyareas of astrophysics. Current
programs execute in-creasingly varied and sophisticated
observations. Ini-tially, these were almost exclusively
observations ofstars and circumstellar phenomena, but recently,
op-tical interferometry extended even to extragalactictargets.
The field is currently driven forward by the ac-tivities in many
research areas, and important scien-tific results are expected in
the near future. Clearly,the main beneficiaries will be stellar
astrophysics and
15
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S. JANKOV
galactic astronomy, in particular the areas of fun-damental
stellar properties, star and planet forma-tion, and all stages of
stellar evolution, but increas-ingly, optical interferometry will
play more impor-tant role in other areas of astrophysics
particularlyexo-planetary and extragalactic research.
Acknowledgements This work has been supportedby the Ministry of
Science and Technological Devel-opment of the Republic of Serbia
(Project No 146007Inverse problems in astrophysics :
Interferometryand Spectrophotometry of stars).
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16
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OPTICAL INTERFEROMETRY I
ASTRONOMSKA OPTIQKA INTERFEROMETRIJA.I METODE I INSTRUMENTI
S. Jankov
Astronomical Observatory Belgrade, Volgina 7, 11060 Belgrade 38,
SerbiaEmail: [email protected]
UDK 520.3613Pregledni rad po pozivu
Prethodna dekada je videla ostva-ree velikih interferometrijskih
projeka-ta ukuqujui teleskope klase 8-10m istometarske
interferometrijske baze. Mo-derni kompjuteri i kontrolna
tehnologijasu omoguili interferometrijsku kombi-naciju svetlosti sa
razdvojenih teleskopakako u vidivom tako i u infracrvenomreimu.
Oslikavae sa rezolucijom od mili-luqne sekunde i astrometrija sa
preciznoxuod mikro-luqne sekunde su tako postale real-nost. Ovde
dajem pregled metoda i instru-mentacije koje odgovaraju tekuem stau
uoblasti astronomske optiqke interferomet-rije. Prvo, ovaj pregled
sumarizuje razvoj odpionirskih radova Fizoa i Majkelsona. Za-tim su
opisane osnovne dostupne veliqine aposle toga sledi diskusija o
osnovnim prin-
cipima konstrukcije modernih interferome-tara. Osnovne
interferometrijske tehnike kaoxto su interferometrija spekl i
aper-turnih maski, sinteza aperture i nu-ling interferometrija su
takoe diskuto-vani. Koristei iskustvo steqeno od prox-lih i
postojeih instrumenata da bih ilus-trovao vane stavke, posebno
razmatram novugeneraciju velikih interferometara koji sunedavno
staveni u upotrebu (posebno QARA,KEK, VLT i LBT interferometri). Na
krajudiskutujem dugoroqnu budunost optiqke in-terferometrije
ukuqujui mogunost novihvelikih projekata za zemaske
opservatorijekao i predviaa za kosmiqku interferomet-riju.
17