NATO ADVANCED STUDY INSTITUTE
OPTICS IN ASTROPHYSICS
ASTRONOMICAL SPECKLE INTERFEROMETRY
Review lecture by Y.Balega from the Special Astrophysical Observatory, Zelentchuk, Russia
Institut d‟etudes scientifiques de Cargese, 2002
2. Speckle phenomena
3. Object-image relation and the Fried parameter
4. Labeyrie’s proposal
5. Speckle interferometry technology: cameras, detectors, photon noise, photon-counting hole.
6. Speckle holography
7. Shift-and-add algorithm
8. Knox-Thomson method
9. Speckle masking
10. Differential speckle interferometry
11. Astronomical applications of speckle interferometry
Everybody knows that telescopes make things bigger: the bigger the telescope the more there is to see. However, atmospheric turbulence limits angular resolution to about 1 arcsec; consequently, a 10-m telescope has approximately the same resolving capabilities as a small 10-cm diameter tube. Resolution limitations caused by the atmospheric smearing, were first overcame by Michelson and his collaborators in 1920s by making use of interferometric properties of light. As you all know, they succeeded to measure the angular diameter of the red supergiant star Betelgeuse using the 5 m baseline installed at the 100-inch telescope. Similar methods have been later used extensively in radio astronomy and in the optical intensity interferometry to reach an angular resolution of 0.001 arcsec and better.
Labeyrie (1970) first pointed out that very short exposure photographs contain information on scales at the telescope diffraction limit. In the 70s the technique called speckle interferometry became one of the most promising developments in the optical observational astronomy. It became an extremely active field scientifically with important contributions made to a wide range of topics in stellar astrophysics. Here
we will make a short review of the speckle interferometry and present the most important results obtained by astronomers using this method at large telescopes.
2. Speckle phenomena
The short-exposure image of a point source formed by a corrugated wave front is composed of numerous short-lived speckles. Speckles result from the interference of
, distributed over the full light from many coherent patches, of typical diameter r0
aperture of the telescope. J.Texereau (1963) was one of the first who described correctly the speckle phenomena in his study of limitations of the image quality in large telescopes. He observed boiling light granules visually using a strong eye-piece and by means of sensitive photographic films in the focus of the 1.93m telescope at Haute-Provence. Earlier, visual observers of binary stars, such as W.Finsen, could use the speckle structure of a binary star image to yield useful information about the separations and position angles between the components. In this manner they have been used the speckle interferometry method without knowing about it. A single r-size subpupil would form a PSF of width ?( / r imposed by diffraction. 00
Diffraction pattern produced by a circular aperture is called the Airy spot. It is made of a very bright central spot surrounded by rings, which are alternately dark and bright. The radius of the first dark ring of the diffraction spot is equal to:
， = 1.22 ( / D
We now consider two small-size subpupils, separated by a distance ? d, pierced into
an opaque screen in front of the telescope aperture. Each small opening diffracts the incoming light, which is spread over the focal plane. The observed phenomena of a modulation pattern of linear interference fringes, is known as Young „s fringes. These are normal to the line joining the two subpupils, and the distance that separates a dark from a bright fringe (fringe width) is ? ( / d.
Atmospheric turbulence causes random phase fluctuations of the incoming optical wave front. As a result of the randomly varying phase difference between the two subpupils, the fringe pattern moves by a random amount. Phase fluctuations are
. In typical equivalent to wave front tilts if the assumed subpupils are smaller than r0
atmospheric conditions, these fluctuations become uncorrelated when the baseline d
exceeds about 10 cm. When the phase shift between the subpupils is equal or greater than the fringe spacing ( / d, fringes will disappear in a long exposure. Instead of making a long exposure, one can follow their motion by recording a sequence of short exposures, short enough to “freeze” the instantaneous fringe pattern.
Now let us introduce a third small subpupil in front of the telescope, which is not collinear with the former two (fig.1). The result will be the pattern of three intersecting moving fringe systems produced by three nonredundant pairs of subpupils. When these fringes interfere, an enhanced bright speckles of width ? ( / d appear.
Fig.1. Speckle pattern formation.
The number of speckles N per image is defined by the ratio of the area occupied by s
the seeing disk ? ( / rto the area of a single speckle 0
222 N = (( / r): (( / D ) =( D / r ) . s0 0
If the atmospheric seeing is good, this corresponds to about 1500 speckles in a single short exposure image taken with an 8-m telescope. More accurate estimates give three times less amount of speckles because of the filling factor k = 0.342.
(Insert Fig.2 “Speckle image from the 6-m telescope”)
is defined by the velocity dispersion in the turbulent The speckle lifetime ：？；
： ? r/ ？； . 0
Under typical atmospheric conditions, speckle boiling can be frozen with exposures in the range 0.01 – 0.05 sec.
3. Object-image relation and the Fried parameter
If the atmospheric degradations are isoplanatic all over the telescope field of view, the irradiance distribution from the object o(x) is related to the long-exposure (ensemble
averaged) observed irradiance distribution < i(x) > by a convolution relation
< i(x) > = o(x) ) < s(x) >
where the point spread function < s(x) > is simply the long-exposure image of a point
In Fourier space it becomes
< I(u) > = O(u) ? < S(u) >
where < S(u) > is the optical transfer function of the system “telescope plus
–1atmosphere”, and u is the spatial frequency vector expressed in radian . For long
exposures, < S(u) > is the product of the transfer function of the telescope T(u) with the atmospheric transfer function equal to the coherence function B(u) (Roddier 1981):
< S(u) > = T(u) ? B(u) .
Similar to the bandwidth in radio electronics, Fried (1966) proposed to define the
resolving power of the telescope as the integral of the optical transfer function:
R = ， < S(u) > du = ， T(u) ? B(u) du .
For a large diameter telescope the resolving power depends only upon turbulence and
R = ， B(u) du.
has a meaning of the critical telescope diameter for which The Fried parameter r0
， B(u) du = ， T(u) du,
Resolving power of the telescope is limited by the telescope only if D； r, in other 0
cases it is limited by the atmosphere. The relation between atmospheric coherence function and r is expressed by: 0
-5/3 B(u) = exp –3.44 ((u / r ) (Fried 1966). 0
Finally, the following two relations show the dependence of rupon the refraction 0
index structure constant Cn (from the Obukhov‟s law: the structure function Dn(，) = 22/3Cn ， (Tatarski 1961)) and the wavelength ( :
2-12 r = [0.423 k (cos z) ， Cn(h) dh, 0
6/5 r ? ( , 0
here k = 2； / ( , and z is the zenith distance.
4. Labeyrie’s method
In the case of short exposures, the image intensity distribution i(x) is again related to
the object distribution o(x) by a convolution relation
i(x) = o(x) ) s(x) .
The Fourier transform of this intensity recording is
? S(u), I(u) = O(u)
222 and ( I(u) ( = (O(u) ( ? (S(u) (,
2 2 here (O(u) ( is the object power spectrum and (S(u) (is the energy spectrum of a
2 point source image. (S(u) (describes how the spectral components of the image
were transmitted by the atmosphere and the telescope. At every moment this function is not known but its time-averaged value can be determined if the seeing conditions do not change. Time averaging means that
222 <( I(u) (> = (O(u) ( ? <(S(u) (>
which leads to
222 (O(u) ( = <( I(u) (> / <(S(u) (>.
Eq. … gives the essence of the Labeyrie‟s (1970) proposal, namely: if the time-
averaged speckle transfer function is determined from the ensemble of its instant recordings, the intensity of the object‟s Fourier transform can be reconstructed (see
2Fig.3). Note that the short-exposure <(S(u) (> is not the same as < S(u) > in eq.(…).
It includes the high spatial frequency component, which extends, as it was shown both theoretically and from observations, up to the diffraction cut-off limit of the telescope.
Fig.3. Labeyrie‟s speckle interferometry.
Analytical expression for the power spectrum of short-exposure images was first proposed by Korff (1973). Its asymptotic behavior is presented in Fig.2, showing that
2> extends up to the ideal telescope diffraction cut-off frequency. That <(S(u) (
means that the typical speckle size is of the order of the Airy pattern of the given
2aperture. The low-frequency part of <(S(u)(> corresponds to a long exposure image
with the wave front tilts compensated.
Fig.3 explains the difference between the long-exposure image acquisition and the Labeyrie‟s reconstruction in the case of close binary star observations. A long-
exposure photograph of an object is equivalent to adding together or averaging many short photographs. The long-exposure transfer function is not diffraction limited and the result is a very blurred image. In speckle interferometry, the turbulent noise is averaged by composition a large number of short-exposure images.
Fig.4. Speckle interferometry transfer function.
It is also convenient to illustrate the Labeyrie‟s speckle technique in the simple case of an equal magnitudes double star with an angular separation between the
( / rsize. In this case the image represents components less than the seeing disk of ? 0
a superposition of two identical speckle patterns; vectors, which connect individual speckles from the two components, are equal to the projected position of the stars on the sky (fig.5). High-angular resolution information cannot be extracted from such
Insert Fig.5. “Binary star image reconstruction by speckle interferometry”
blurred images visually, however the binary star can be studied from its Fourier transform pattern or from its averaged autocorrelation
< ， i (x) i (x + x? ) dx >
which is the inverse Fourier transform of the image power spectrum.
An important consideration for narrow-angle speckle interferometry is the sky coverage defined by the isoplanatic patch (see the lecture of C.Paterson).
5. Speckle interferometry technology: cameras, detectors, photon noise,
The ideal speckle camera should satisfy the following two conditions:
- high-resolution short-exposure image recording;
- spectral bandwidth selection.
Usually it consists of the following components (fig.6):
- shutter for cutting out short exposures in the range 0.01 – 0.1 s:
- microscope objectives set to match the speckle size to the pixel size of the
- filter wheel or a diffraction grating monochromator for bandwidth
- atmospheric dispersion compensation prism;