Sky demonstration of potential for ground layer adaptive optics correction

Transcription

1 Sky demonstration of potential for ground layer adaptive optics correction Christoph J. Baranec, Michael Lloyd-Hart, Johanan L. Codona, N. Mark Milton Center for Astronomical Adaptive Optics, Steward Observatory, 933 N. Cherry Ave., Tucson, AZ ABSTRACT Observations have been made at the Steward Observatory 1.55 m telescope of a four-star asterism in the constellation Serpens Cauda, using a Shack-Hartmann wavefront sensor. The stars are all within a 2 arcminute field, and range in apparent brightness from m v of 9.4 to 1.6. The instrument placed a 5 5 array of square subapertures across the pupil of the telescope, and had sufficient field of view to allow wavefront data to be recorded from all four stars simultaneously. Snapshots at 1/3 s exposure time were recorded, with no temporal coherence between exposures. We have reconstructed the first 2 Zernike modes from the slope data for each star. In a preliminary analysis, we show that the wavefront aberration in each star can be roughly halved by subtracting the average of the wavefronts from the other three stars. The averages represent estimates of the aberration introduced by the lowest few hundred meters of the atmosphere, so the result provides an early indication of the potential for image sharpening by compensation of boundary layer turbulence. Keywords: Wavefront sensors, adaptive optics, ground layer adaptive optics 1. INTRODUCTION Ground layer adaptive optics (GLAO) correction, first suggested by Rigaut 1, is a method for correcting the wavefront errors caused by turbulence close to the telescope. By averaging the wavefronts of multiple guide stars over a wide field of view, one can calculate an estimate of the errors introduced by the ground layer. Compensating this error with a deformable mirror conjugated close to the telescope s entrance pupil will partially correct the wavefronts from all objects in the field of view simultaneously. Numerical simulations 1,2,3 suggest that, when correctly implemented, this can be an effective technique for achieving modest correction over fields of view of perhaps 1 to 2 arcminutes. GLAO has yet to be demonstrated at any telescope, and its potential is very uncertain because the relative contributions to the wavefront aberration from turbulent boundary layers and those at high altitude are not generally well understood. The surest way to explore the potential is to measure explicitly and simultaneously the wavefront errors in several stars within some field, and to calculate the degree to which one star s wavefront can be estimated on the basis of measurements from the other stars. Ragazzoni, Marchetti, and Valente first attempted this at the Telescopio Nazionale Galilei 4, using the simple expedient of defocusing their camera to allow wavefront estimates to be computed. Here, we report the first effort to record multiple stellar wavefronts simultaneously using a specifically designed wavefront sensor, and preliminary results derived from the data. 2. EXPERIMENTAL DESIGN The goal of the experiment was to explore the feasibility of ground layer adaptive optics correction. For this purpose, the field of view for finding reference stars was restricted to 3 arcminutes. While this is much smaller than the field that simulations suggest may benefit from GLAO, it greatly eases constraints on the optical design of the wavefront sensor. To carry out this experiment, we then needed to find a collection of reasonably bright stars, m v > 11, within a ~3 arcminute field of view. We found such a set of stars in the constellation Serpens Cauda with visual magnitudes ranging from 9.4 to 1.6 with separations from a central star of 56 to 76 arcseconds. Figure 1 shows the relevant section of a DSS image.

2 Figure 1. Sky picture of target stars. Taken from DSS2/J/POSSII 2.1. Specifications and design of the WFS The main challenge in this experiment was to design an appropriate camera that could capture wavefront information from the multiple stars in the selected field simultaneously. We chose to design a wavefront sensor that could image multiple Shack-Hartmann patterns onto a single CCD chip. The camera was designed to cover a 2.5 arcminute square field of view with a minimum separation of 3 arcseconds between each star so the patterns would not overlap. The pupil was broken up into 5 by 5 subapertures, projecting to 31 cm squares on the primary mirror, of which 2 were illuminated, allowing us to sense 4 degrees of freedom. The final plate scale on the camera is 1.68 pixels per arcsecond. Figure 2. Zemax layout of WFS design showing simulated fields of target stars. From left to right, the optical components: Focal plane of telescope (1), where images of stars are formed, collimating lens (2) which images the pupil onto the lenslet array (3). Next is a field lens (4) that corrects aberrations and following is a relay system (5,6) which images the Shack-Hartmann spot patterns onto a CCD chip (7).

3 2.2. Implementation on the telescope The camera was designed to be used at the f/13.5 Cassegrain focus of Steward Observatory s 1.55 m Kuiper Telescope. The plate scale at best focus is 11 microns per arcsecond with a useful field of view diameter of at least 435 arcseconds. Typical seeing at the site, Mt. Bigelow, just north of Tucson, Arizona, is 1-2 arcseconds. Figure 3. The assembled wavefront sensor camera attached to the Cassegrain focus of the 1.55 m Kuiper Telescope. 3. DATA COLLECTION Observations using the wavefront sensor camera were done over the nights of June 12 th - 16 th, 3. The data used for analysis was taken from 1am to 3am on the morning of the 17th. During that time, a total of 625 frames of data from the sky in batches of 25 frames. After each set, the telescope was offset to the sky and 2 background frames were recorded, for a total of 5 background frames. The sky background was found to remain constant over the night so the mean of all the sky frames was used for background subtraction in the data analysis. Figure 4 shows an example of a single background subtracted data frame. Each frame had an exposure time of 31 milliseconds. There was a delay of ~2 seconds between each exposure, which means that successive frames were temporally uncorrelated.

4 Figure 4. An example data frame after background subtraction, showing the Shack-Hartmann patterns from the four stars in the selected field. The field shown is 2.5 arcminutes square. 4. DATA ANALYSIS AND WAVEFRONT RECONSTRUCTION 4.1. Data reduction / Centroid offset vector calculation In order to recover wavefront information from the data frames, the spots from each subaperture in the Shack- Hartmann patterns had to be accurately located. Using a correlation technique, each frame was convolved with a 2-D Gaussian of width 2 pixels. The coordinates of the spots were taken to be the peaks in the correlation. The zero offset position for the spot in each subaperture was assumed to be the average position of the corresponding batch of 25 frames. The zero offsets were then used to quantify the distortion seen in figure 4, and to compute an adjustment to the plate scale as a function of position on the camera. Finally, centroid offset vectors were calculated from the spot positions frame by frame by subtracting the zero offsets, and applying a correction for the measured distortion Ground layer estimation using a single subaperture The first analysis of the data was to estimate the ground layer contribution for a single subaperture. By measuring the centroid offset vectors for a single subaperture, we are effectively measuring the tip and tilt over a 31 cm square telescope. Given the separations of the stars, the overlap of the 31 cm beams from the central and outlying stars extends only to a height of ~1 m. Centroid data for the same subaperture in each of the three outer stars were averaged to give an estimate of the ground layer tilt, and subtracted from the centroid data for the central star. Figure 5 shows a comparison of the central star s centroid offset vector with and without subtraction of the ground layer

5 estimate for the first 91 frames of data. The RMS length of the offset vector has been reduced by a factor of 2.28 from.567 arcseconds in the uncorrected case to.249 arcseconds in the corrected case. Note also that the roughly circular distribution of the offset vectors around zero leads us to believe that the wavefront tilt errors are due largely to turbulence in the atmosphere rather than telescope wobble. Centroid Offset Positions - Frames 1-91 Subap (2,2) centroid offsets corrected centroid offsets 5 Y (.1 arcseconds ) X (.1 arcseconds ) Figure 5. Scatter plot of the centroid vector from a single 31-cm subaperture for the central star, with and without subtraction of the average tip-tilt computed from the three outer stars Ground layer estimation of full aperture using centroid offset vectors In an extension to the above analysis, the same type of correction was applied to all the subapertures instead of just one. Again we averaged the centroid offset vectors for each subaperture of the three outer stars and used this as a correction to the offsets, but this time we applied the correction to each of the stars in the field. We can see from table 1 the improvement in RMS centroid offset vector length over the entire aperture for each star. In this simplistic ground layer correction we can see that there is significant improvement in the wavefront slopes for each of the stars observed. Star Upper Left Central Right Lower Right Uncorrected Centroid Offset Vector (arcseconds) Corrected Centroid Offset Vector (arcseconds) Table 1. RMS centroid vector length over the entire aperture with and without subtraction of the ground layer estimate.

6 4.4. Wavefront reconstruction The next step was to calculate estimates of the wavefront, expressed as sets of low-order Zernike polynomial coefficients, using a reconstructor matrix derived from a numerical model of the imaging system. With just 2 illuminated subapertures, we restricted ourselves to 2 Zernike modes, so the reconstruction calculation remained slightly overdetermined. Figure 6 shows an example of the reconstructed optical path differences superimposed on the original sky image of our selected stars. Figure 6. Sky image with superimposed reconstructed phase maps from a single data frame. To compute a ground layer estimate, we averaged the Zernike polynomial coefficients of the three outer stars. These coefficients were then subtracted from the coefficients calculated for each individual wavefront. Figure 7 shows the calculated RMS Zernike polynomial coefficients in both the uncorrected and corrected cases. From the uncorrected data, we estimate that the seeing parameter r was about 11 cm at 5 nm. Over the four stars in the field there is modest improvement in wavefront error for all modes and all stars, with a total compensation for all 2 modes of about a factor of 2, from 818 nm to 445 nm. Upper Left Star Right Star Lower Right Star Central Star Zernike Order rms wavefront error rms wavefront error % Correction % Correction % Correction % Correction without correction (nm) with correction (nm) Total: Table 2. Correction by Zernike radial order in each star by subtraction of the mean wavefront measured from the three outlier stars.

7 8 7 RMS Zernike Polynomial Coefficients Over 11 Frames: Upper Left Star Uncorrected After Ground Layer Correction 8 7 RMS Zernike Polynomial Coefficients Over 11 Frames: Central Star Uncorrected After Ground Layer Correction 6 6 Coefficient in Nanometers Coefficient in Nanometers Zernike Mode Number Zernike Mode Number 8 7 RMS Zernike Polynomial Coefficients Over 11 Frames: Right Star Uncorrected After Ground Layer Correction RMS Zernike Polynomial Coefficients Over 11 Frames: Lower Right Star 8 Uncorrected After Ground Layer Correction Coefficient in Nanometers Coefficient in Nanometers Zernike Mode Number Zernike Mode Number Mode # Order Name Mode # Order Name 1 1 Tilt x 11 4 Secondary astigmatism y 2 1 Tilt y 12 4 Spherical 3 2 Astigmatism x 13 4 Secondary astigmatism x 4 2 Power 14 4 Tetrafoil x 5 2 Astigmatism y 15 5 Pentafoil y 6 3 Trefoil y 16 5 Secondary tetrafoil y 7 3 Coma y 17 5 Secondary coma y 8 3 Coma x 18 5 Secondary coma x 9 3 Trefoil x 19 5 Secondary tetrafoil x 1 4 Tetrafoil y 2 5 Pentafoil x Figure 7. Uncorrected and ground layer corrected RMS Zernike polynomial coefficients. 5. CONCLUSIONS AND FUTURE WORK Simulations have suggested that GLAO will improve rms wavefront error by about a factor of two, and this is borne out by the present results. One major limitation to this work is that we have no independent knowledge of the relative strength of the ground layer compared to other turbulence layers in the atmosphere. Despite this lack, an estimate can be obtained of the mean height of non-ground layer aberration from the angular decorrelation of the Zernike modes as a function of their radial order. We are modeling this now. A tomographic reconstruction algorithm will be used to estimate contributions from two layers assumed to be at this mean height, and the ground. This would represent the first validation of tomography with sky data, and if successful, will improve the estimates of the central star s wavefront.

8 In further experiments at the 1.55 m, the present camera will be replaced with one having a faster readout and lower read noise, allowing some degree of temporal coherence to be maintained between exposures. An upgraded optical system is also being designed to allow the wavefront sensor to be used at the 6.5 m MMT Observatory on Mt. Hopkins. There we will utilize another local 1.2 m telescope to obtain simultaneous measurements of the C n 2 profile via SCIDAR. From these data, we expect to measure the average fractional power in the ground layer aberration, and so to quantify the effectiveness of GLAO. We will also be able to anchor tomographic solutions for the 3-d aberration structure against the known turbulence profile. 6. ACKNOWLEDGEMENTS We are grateful for assistance from James Georges, Tom Stalcup, and Matt Rademacher. This work has been supported by the National Science Foundation under grant AST REFERENCES 1. F. Rigaut, Ground-Conjugate Wide Field Adaptive Optics for the ELTs, Proc. ESO Beyond Conventional Adaptive Optics, Venice, Mark R. Chun, Gains from a ground only adaptive optics system, Proc. SPIE Conference on Adaptive Optical System Technologies II, ed. Peter L. Wizinowich & D. Bonaccini, 4839, 94-98, Kona, HI, M. Lloyd-Hart, F. Wildi, G. Brusa, Lessons learned from the first adaptive secondary mirror, Proc. SPIE Conference on Astronomical Adaptive Optics Systems and Applications, ed. Robert K Tyson & M. Lloyd-Hart, 5169, San Diego, CA, R. Ragazzoni, E Marchetti, & G. Valente, Adaptive-optics corrections available for the whole sky, Nature, 43, 54-56,.