Before contemplating an optical interferometer on the Moon one must first review the accomplishments achieved by this technology in scientific applications for astronomy. This will be done by presenting the technical status of optical interferometry as achieved by the Mount Wilson Optical Interferometer. The further developments needed for a future lunar-based interferometer are discussed

Colavita, M.

Johnston, Ken J.

Mozurkewich, D.

Shao, Michael

Simon, R. S.

English

NASA Technical Reports Server

N93-13582
THE MOUNT WILSON OPTICAL INTERFEROMETER:
THE FIRST AUTOMATED INSTRUMENT
AND THE PROSPECTS FOR LUNAR INTERFEROMETRY
ILJ. Johnston, D. Mozurkewich, R.S. Simon
E.O. HuIburt Center for Space Science
Naval Research Laboratory
Washington, DC 20375
M. Shao, M. Colavita
Smithsonian Astrophysical Observatory
Cambrige, Massachusetts
Iam tam 
Before contemplating an optical interferometer on the Moon one must first review the
accomplishments achieved by this technology in scientific applications for astronomy. This will
be done by presenting the technical status of optical interferometry as achieved by the Mount
Wilson Optical Interferometer. The further developments needed for a future lunar-based
interferometer will be discussed.
Ihakgz and
Long Baseline Optical Interferometry (LBOI) is the use of discrete elements to obtain the
detailed spatial structure of celestial objects. The light received from two independent apertures is
brought together with the light paths being made equal through the use of a delay line or correcting
plate as shown in figure 1. The first known successful application of LBOI to measure the
: diameters of stars was accomplished by A. Michelson in 1920 with the successful measurement of
the diameter of Betelgeuse. This was done using mirrors mounted on a 20-ft beam placed on the
100-in. Mount Wilson telescope. The diameters of six stars, all giants and supergiants of late
spectral class, were measured. This interferometer was abandoned because it was too difficult to
stabilize the light paths. In 1960, R. Hanbury Brown developed the technique of intensity
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https://ntrs.nasa.gov/search.jsp?R=19930004394 2020-03-17T10:19:45+00:00Z
interferometrywith whichthediametersof32brightbluestars were measured. In the 1970s,
Labeyrie directly combined the light beams from two telescopes. He obtained fringes on Vega with
a 12-m baseline using the I2T interferometer. These instruments all used the visibility of the
crosscorre!ated signal to determine the diameters of stars. The crosscorrelated signal is a
complex number containing not only the visibility or amplitude but also phase. The phase gives
detailed positional information on the source of the signal.
In 1968, R. Hanbury Brown stated that Michelson's interferometer, or one in which the
signals are mixed before detection, is applicable to many astronomical problems but technical
problems remained. To produce significant results, the separation must be considerably larger
than the 20-ft baseline originally used by Michelson, the instrument must be free from errors due
to atmospheric seeing, and the results must be recorded in some objective way which is
independent of the skill of the individual observers.
The sensitivity of an interferometer is proportional to the optical bandpass, the area of the
collecting aperture, and the length of the coherence integration time. The diameter of the
collecting aperture is limited to approximately 10 cm for one arcsec seeing by the spatial coherence
of the Earth's atmosphere at optical wavelengths (5500 A). That is to say that the randomness of the
turbulence in the atmosphere leads to randomization of the phase on a single aperture or mirror
over a spatial distance greater than 10 cm. Similarly the coherent integration time is also limited
to approximately 10. To maximize the sensitivity, wide bandpasses are needed. This in turn
results in an interference pattern which is very narrow. This pattern is only a few fringes wide
and requires active fringe tracking to maintain coherence in the presence of atmospheric
turbulence. This in turn limits the sensitivity of Earth-based interferometers to approximately 10
magnitude (Visual) when only a single atmospheric coherence cell is used.
The Mount Wilson Interferometer
As noted, before 1978 alI LBOIused only the amplitude of the crosscorrelated signal. The
first phase coherent optical interferometer that recorded the amplitude and phase of the
crosscorrelated signal using a phototube, computer, and precise delay li_ne to track the differential
atmosphere phase path was developed by Shao and Staelin, who tracked the "white light" fringe of
Polaris in 1979. This instrument has been improved by the use of the larger apertures, longer
baselines, and wider sky coverage into the Mark II and Mark III interferometers.
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Figure2 shows a schematicdiagram ofan opticalinterferometeras embodied in the Mark
I-IIIinterferometers.The lightfrom a celestialobjectinthe two arms ofthe interferometeris
brought togetheratthe photomultiplierwhere,when fringesare observed,thereisa maximum in
intensity.The major innovationisthe ditherdelayline,which vibratesat a frequencyofa kHz
with an amplitudeofabout a wavelength ofthe lightbeing observed.Once fringesare detected,the
ditherdelaylineallowsthem tobe trackedatthe ms time scale.This isaccomplishedalsoby the
precisiondelaylinethatcan be setto 100 A accuratelyand readout to50 A accuracy.By varying
the delayline,the peak amplitudeofthe fringepatterncan be scanned,thus compensating for
variablepath lengthdelays.A computer controlsthe positionofthe delayline,findsthe maximum
in the fringeamplitude,and recordsthe delay,amplitude,and phase ofthe crosscorrelatedsignal.
Thus the path lengthsin thetwo arms oftheinterferometercan be made equal and trackedon
timescalesofms. Sincethe timescaleforturbulencein the Earth'satmosphere isoforder 10 m, the
interferometercan compensate forphase fluctuationsdue to atmosphericturbulence.
In addition to measuring the sizes of stars, interferometry can also precisely measure
their positions. The USNO has an active program in astrometry and was seeking new technology
to improve the accuracy of the positions of stars. Therefore the development of optical/IR
interferometry for astrometry was sponsored by the Office of Naval Research. In 1982, a joint
program in optical interferometry aimed at astrometry was undertaken by the Naval Research
Laboratory (NRL), the United States Naval Observatory, the Smithsonian Astrophysical
Observatory (SAO), and the Massachusetts Institute of Technology (MIT). The result was the
Mark III interferometer.
Since the major impetus of the NRL/USNO/SAO/MIT program was primarily astrometry
or the precise measurement of the positions of stars, the apertures collecting the light from the stars
had to be made to rapidly and automatically switch from star to star in a preprogrammed
sequence. This required, aside from pointing the mirrors precisely at the stars, that the finding
and tracking of the central fringe also be totally automated.
The Mark III stellar interferometer was built to demonstrate fundamental astrometric
measurements. A secondary goal was to initiate a program of accurate stellar diameter
measurements leading to an instrument for imaging stars. Since astrometric measurements of
high precision require repeated measurements, among the goals of this instrument were that it be
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easilyoperable,reliableandcapable of extremely accurate measurements. A number of active
subsystems were incorporated into the instrument to achieve these goals. The interferometer can
be divided into five major subsystems: (1) a star tracker, (2) the optical delay line, (3) the stellar
fringe acquisition and tracking system, (4) a laser metrology system, and (5) the siderostat
pointing and control system.
Figure 3 shows the Mark III stellar interferometer, which is locaterd on Mt. Wilson,
approximately 80 m east of the 100 in. telescope where Michelson and Pease made the first
measurements of stellar diameters. Six and ten meters north, south and east of the central
building are located the siderostat piers. In the figure, the siderostats are mounted on the
innermost piers. The siderostats are located in huts that provide weather protection for the
siderostats. The hinged roofs of the huts swing open for observing. Local seeing effects are
reduced by routing the starlight through vacuum pipes from the siderostats into the main building.
The main building contains the optics that combine the beams, the delay line, etc. The light is
directed into the main building by the flats and is directed by piezoelectrically controlled mirrors
toward the vacuum delay line. After reflection by the delay line's retroreflectors, the beams are
combined at the beamsplitter. Part of the light is directed into the optical fibers that feed the
phototubes and part of it is directed toward the star trackers. The trailer to the let_ of the main
building houses the computers and observers. There are four possible baseline configurations
form 8.3 m NE-SW to 20 m N-S as shown in figure 4. Also shown is a 4-30 m variable baseline
along a N-S line that is capable of amplitude measurements that can be used to measure stellar
diameters and evaluate atmospheric turbulence.
Astrometric measurements have been carried out with this instrument. As already stated,
these observations involve observing as large a number of stars as possible to measure the baseline
length and solve for the star positions. The number of stars observed is limited by how quickly the
siderostats can move from star to star and the integration time of the individual measurements.
At the present time it takes about 1 min to move between stars and 1 min to obtain an ample amount
of data on the individual star. Figure 4 shows the observed delays as a function of time for the
observations made on November 11, 1986, with a 12 m N-S baseline. The delay changes as the
orientation of the baseline to the stars varies with the Earth's rotation. Rapid measurements must
be made among several stars to determine the baseline length and its orientation as the baseline
length drifts by approximately a micron an hour.
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Measurements made in the fall of 1986 demonstrated that one-color observations could
determine star positions in one coordinate to 20 mas. Measurements made in 1988 in two colors
made during six nights of observation using the 12 N-S and 8.3 S-E m baselines of 12 stars
displayed a formal accuracy of 6 to 10 mas in both celestial coordinates. It is very difficult to
compare the accuracies of these star positions to anything available at the present time. The
formal accuracy of the best star positions, i.e., the FK5 catalog, which is a compilation of the best
available optical data, is at the 50 mas level at epoch 1988. The positions of these stars were in
agreement with the FK5 positions at the 50 mas level for over half the stars. Repeated
measurements with the Mt. Wilson instrument will have to be made to ascertain the systematic
errors on the optical interferometric measurements.
Stellar diameters have also been measured with this instrument. A 12 m baseline at
optical wavelengths has a minimum fringe spacing of 8 mas on the sky. If we consider stars to be
spherical, which is a rough first approximation which is probably correct at the few percent level,
then a one-dimensional variable baseline can be used to measure stellar diameters, the variable
4-30 m baseline has measured the sizes of about 20 stars as of November 1988. Figure 6 shows the
observations for the star alpha Arietes. The least squares fit to the stellar diameter is 6.29 mas
with an RMS error 0.12 mas. The dotted lines in figure 6 show the visibility curves for a uniform
cylinder having a diameter 0.25 mas larger and smaller than the least squares value.
Observation with the Mt. Wilson interferometer will continue in 1989 to evaluate the effects
of the atmosphere over long baselines, extend the measurements of stellar diameters and to repeat
the astrometric measurements. Further, this instrument will be used to demonstrate prototype
systems for future Earth- and space-based optical interferometers.
Develovments °Necessary_ for Ima_ng
The Mt. Wilson interferometer is a two-element interferometer. It has shown that
interferometric fringes can be obtained over baselines of length 5 to 30 m and that the operation of
an interferometer can be automated from pointing the telescopes to setting the delay line and
phasing the instrument.
For true interferometric imaging of objects at optical wavelengths, three key technologies
must be developed: simultaneous combination of beams from multiple elements, longer optical
delay lines, and a metrology system to precisely monitor the geometry of the instrument. Longer
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delaylines are needed to compensate for the variable arrival of the signals at the apertures if
baselines longer than 20 m are to be used. These necessary developments should first be
undertaken and proven with a ground-based instrument.
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Thus, the next logical step in the development of optical interferometry is to build a ground-
based instrument for high angular resolution imaging of stars, stellar systems, and other
celestial objects. This instrument will be the highest resolution imaging instrument ever used at
optical wavelengths, and will achieve resolutions exceeding even those available from Very Long
Baseline Interferometry techniques in the radio in all but the shortest radio wavelengths. This
instrument will represent a tremendous advance over any currently existing optical
interferometer, offering improvements of a factor of 10 or more in sensitivity, resolution, and
imaging speed.
The instrument will be located on a suitable mountaintop where atmospheric conditions
will facilitate optimum performance. The overall size will be about 250 in diameter, with at least
six independent telescopes or siderostats forming a reconfigurable array. Light from these
telescopes will be combined interferometrically in a central optics laboratory to produce the basic
visibility amplitude and phase data used to form images.
The research carried out with such an interferometer will have a profound effect on the
technology of imaging objects at great distances and will greatly aid in our understanding of the
physical attributes such as size, shape, distance, and mass of celestial objects such as stars. The
capability of imaging objects will be at least two orders of magnitude better than the Hubble Space
Telescope. This improvement in resolution will allow many important astronomical discoveries
to be made.
This interferometer is to be the first to image objects directly at optical wavelengths using
both the amplitude and phase of the crosscorrelated signal. Early in the program the technology of
simultaneously combining three beams and making phase closure measurements should be
accomplished. Later the multiple-element instrument will demonstrate from the ground the
capabilities of imaging celestial objects. After this an evaluation of the capabilities of using this
technology in a space-based system can be made. The first space-based interferometer will be in
Earth orbit but later instruments could be located on the Moon.
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Interferometrvin Snace
The advantage of space-based optical interferometry is freedom from viewing objects
through the Earth's atmosphere. For an Earth-based optical interferometer, the sensitivity of the
instruments is proportional to seeing to the third power. This allows the phase across the observing
aperture to be constant whereas on the Earth, for seeing of order one arcsec, apertures larger than
10 cm cannot be used for a single coherent aperture at 5500/_ Apertures of size 10 cm must be
summed. Further, integration times greater than 10 mas cannot be used (again assuming one
arcsec seeing).
Imaging objects on the Earth with a conventional phase-stable interferometer using a
single atmospheric coherence cell will be limited to objects brighter than 10th magnitude. If by
summing different cells in the Earth's atmosphere increases the collecting area of the
interferometer, or if a natural or artificial reference source can be found or generated to increase
the integration time, fainter magnitudes can be reached. An interferometer in space will not have
this limitation because very large optics can be used. The limitation on the imaging magnitude
will be the stability of the space-based platform and the figure and stability of the large apertures.
Therefore, for a space-based system, the intrinsic stability of the interferometer or the metrology
system, i.e., the system measuring the spatial stability of the interferometer, will set the limiting
magnitude constraint.
For free-flying interferometers, the stability of structures in space must be studied and
metrology systems must be developed to overcome the shortcomings of the structures. The sizes of
stars vary from 60 to 0.1 mas while the size scales of extragalactic objects are of order a few mas or
less. This leads to useful baselines at optical wavelengths of order 10 m (10 mas) to 1000 m
(0.1 mas) for imaging objects. It is very doubtful that structures as large as 100 m would be built for
a free-flying interferometer, thus, for the longer baseline, either multiple free-flying satellites
will be used, or the interferometer will be placed on the Moon.
Other reasonsforplacingan interferometerin space are to improve the instrumental
stabilityand astrometricmeasurements. A free-flyingspace astrometricinstrument would
probablyuse a very shortbaseline,but therewould be no limiton baselinelengthfora lunar
instrumentexceptthatthe objectnot be resolved.For an interferometerbased on the lunar surface,
the problems ofinstrumentalstabilityshouldbe simplerbecause massive structurescan be used
forstability.Note thatthe Mt. Wilson interferometerhas massive pierswhich move atabout a
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micron per hour. Thermal effects should have a much slower timeconstant for a lunar instrument
since thermal effects will vary with the lunar rotation period whereas a sa_llite in earth orbit will
have a much shorter thermal period unless it is sun synchronous.
The Mt. Wilson interferometer has demonstrated the necessary technologies of telescope
coalignment and optical delay lines. Three further developments in ground-based
interferometry are needed before the development of space-based interferometry. These are the
simultaneous combination of beams from multiple elements, longer optical delay lines, and a
metrology systems to precisely monitor the geometry of the instrument. After these developments
take place, the study of space'based structures should be undertaken to evaluate the applicability of
Earth-orbiting interferometers. In the long term, optical interferometry will need baseline
lengths of order 100 m or greater for imaging and multiple free-flying individual apertures, or a
lunar based interferometer must be considered. The technology for a lunar interferometer
appears to be close in hand as it is just an extension of the Mt. Wilson interferometer and should
require a minimum of technology development. Thus the prospect of a lunar interferometer
should be carefully considered when the development of the Moon for scientific purposes is
undertaken.
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The Mount Wilson optical interferometer: The first automated instrument and the prospects for lunar interferometry

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