The Astronomical Journal, 136:312–322, 2008 July c 2008. The American Astronomical Society. All rights reserved. Printed in the U.S.A.
doi:10.1088/0004-6256/136/1/312
CHARGE-COUPLED DEVICE SPECKLE OBSERVATIONS OF BINARY STARS WITH THE WIYN∗ TELESCOPE. V. MEASURES DURING 2001–2006 Elliott P. Horch1,4,5 , William F. van Altena2 , William M. Cyr, Jr3,4,6 , Lori Kinsman-Smith3,4 , Amit Srivastava3,4,7 , and Jing Zhou3,4 1
3
Department of Physics, Southern Connecticut State University, 501 Crescent Street, New Haven, CT 06515, USA;
[email protected] 2 Department of Astronomy, Yale University P.O. Box 208101, New Haven, CT 06520, USA;
[email protected] Department of Physics, University of Massachusetts Dartmouth, 285 Old Westport Road, N. Dartmouth, MA 02747, USA;
[email protected],
[email protected],
[email protected], and
[email protected] Received 2007 December 21; accepted 2008 April 8; published 2008 June 11
ABSTRACT A total of 1067 speckle observations of 345 binary stars are presented. Of these, 161 are double stars first resolved by Hipparcos, 17 are resolved for the first time in the observations presented here, and 21 are stars previously discovered by our program and reported in earlier papers in the series. In 947 cases, a magnitude difference is reported along with the relative astrometry. When comparing to systems with very well-known orbits, we find that the root mean square (rms) deviation in separation residuals is 2.81 ± 0.28 mas, and the rms deviation in position angle residuals is 0.88 ± 0.07◦ . The magnitude difference measures show no significant deviation from Hipparcos photometry, and have average standard deviation of approximately 0.10 mag as judged from repeat observations. Five important systems discovered by Hipparcos are discussed. Key words: astrometry – binaries: visual – techniques: interferometric – techniques: photometric Online-only material: machine-readable and VO tables
step-and-expose pattern to use the full area of a large-format CCD camera for collecting speckle images. Since the inception of RYTSI, the observational program has focused on two main groups of binaries. The first is a list of approximately 100 systems where high-quality orbits already exist, but differential photometry is either nonexistent or in need of improvement before comparisons with theoretical models could be made. The second group of objects has been the Hipparcos Double Star (HDS) discoveries, prioritized by distance. Since the distances to many of these systems are already relatively well known, determining the visual orbital elements for those systems with reasonable periods is the next step en route to determination of the mass sum. Thus, these systems represent many potential tests of stellar evolution theory, if they can be observed with high precision over a period of years. In this regard, some 19 HDSs already have orbit determinations from the work of Balega and his collaborators (Balega et al. 2002, 2005, 2006), S¨oderhjelm (1999), Docobo et al. (2007), and Romero (2007). In this paper, we present results of observations obtained with RYTSI during 2001–2006 at WIYN and discuss five of the more interesting HDS systems.
1. INTRODUCTION Binary stars continue to be a fundamental tool in understanding stellar masses. Correlating mass determinations with other important parameters such as radius, luminosity, and effective temperature gives a basic check of stellar structure and evolution theory. In the past, speckle interferometry provided precise relative astrometry for many sub-arcsecond separation binary systems, but lacked sufficient photometric precision to make meaningful comparisons between masses obtained from subsequent orbit determinations and component magnitudes and colors. The speckle observing program at the WIYN 3.5 m Telescope has been designed to overcome this problem by providing reliable relative photometry of the components of binary systems in addition to high-precision relative astrometry. Observations are taken with a large-format CCD, where in the initial years of the program (1997–2000), the high frame rate necessary for successful speckle observations was handled by row-shifting accumulated charge in a timed pattern that allowed a number of speckle patterns to be accumulated prior to readout (this method is discussed in Horch et al. 1999). In 2001, we began observing with the RIT-Yale Tip-tilt Speckle Imager (RYTSI), which is described in Meyer et al. (2006). RYTSI has a two-axis galvanometer scanning mirror system which executes a timed
2. OBSERVATIONS AND DATA REDUCTION For all observations discussed in this paper, the RYTSI optics package and the imaging CCD were mounted at the Nasmyth port (the so-called “WIYN port”) of the WIYN Telescope. The plate scale delivered by the telescope optics is 9.372 arcsec mm−1 , which is then magnified by the RYTSI so that near-Nyquist sampling of speckles is obtained on the pixels of the CCD detector. Because the telescope has an altitude– azimuth mount, the WIYN port is equipped with an instrument rotator as a part of the instrument assembly subsystem that keeps the orientation of images on the detector constant regardless of sky position. Measurements by WIYN personnel indicate that the uncertainty of orientation angle relative to sky coordinates is less than 0.1◦ for sky positions that are not within a few degrees
∗ The WIYN Observatory is a joint facility of the University of Wisconsin–Madison, Indiana University, Yale University, and the National Optical Astronomy Observatories. 4 Visiting Astronomer, Kitt Peak National Observatory, National Optical Astronomy Observatory, which is operated by the Association of Universities for Research in Astronomy, Inc. (AURA) under cooperative agreement with the National Science Foundation. 5 Adjunct Astronomer, Lowell Observatory. 6 Present Address: Science Department, Norwood High School, 245 Nichols Street, Norwood, MA 02062, USA. 7 Present Address: Bhaktivedanta Institute, ISKCON, Juhu Road, Juhu, Mumbai-4000049, India.
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of the zenith; it is not practical to observe in that region in any case due to the tracking limitations of the telescope’s azimuth drive. We therefore assume that the error introduced in position angle due to instrument rotation is less than 0.1◦ . On all but one observing run described below, a Kodak KAF-4200 CCD was used in combination with RYTSI to obtain the speckle observations. This is a front-illuminated device with 9 µm square pixels in a 2033 × 2048 format. The read noise is approximately ten electrons, and the peak quantum efficiency is 40%. The RYTSI optics were then configured with a collimating lens of focal length 60 mm and a re-imaging lens of focal length 180 mm, leading to a nominal 3× magnification over that provided by the WIYN port. The alternate CCD used in the 2004 February run was the Kitt Peak Mini-Mosaic camera. This device consists of two 2048 × 4096 CCDs positioned side by side in the same instrument dewar. Here, the (square) pixel size is 15 µm, the read noise is approximately six electrons, and the peak quantum efficiency is 80%. Because of the larger pixel size, a 40 mm collimating lens was used in combination with the 180 mm re-imaging lens in RYTSI, giving a nominal magnification of 4.5× over that of the WIYN port. The normal speckle observation when using the Kodak CCD was to configure the scanning mirrors in RYTSI to execute a 16 × 16 grid of 50 ms exposures, with each exposure separated from its nearest neighbors by approximately 2.5 arcsec. This can be collected in 12.8 s, followed by a full-frame readout of the chip. (The electronics for our chip complete the full-frame read in roughly 8 s.) Four such exposures were generally obtained sequentially and stored in the same raw data file as a FITS stack of dimensions 2048×2048×4. Therefore, a typical observation consists of 1024 individual speckle images, and because the CCD’s active area is 2033×2048, each file contains an overscan region which can be used to monitor and remove any drift in the bias level from observation to observation. This is an important feature, especially for obtaining accurate relative photometry, as discussed in Tyler et al. (2007). In the case of Mini-Mosaic speckle files, the mirrors were configured to give a grid pattern of 22 × 22 speckle exposures of 50 ms each and generally only one image was taken giving 484 individual speckle patterns on the target. Due to the slower readout of the Mini-Mosaic camera versus the Kodak camera, it was much more efficient to point to the subsequent target rather than waiting for multiple readouts of the CCDs. This was also not a detriment to the data collected because the increase in the quantum efficiency obtained by using the Mini-Mosaic more than compensated for fewer speckle images in the data file. The normal nightly observing sequence was to begin by collecting data on one or more binaries that were within a few degrees of one another on the sky. Included in the grouping would be the observation of a bright unresolved star chosen from The Bright Star Catalogue (Hoffleit and Jaschek 1982). We would then move to another binary or group of binaries, and include in that grouping another bright point source, and so on. We would therefore observe perhaps 80 to 100 science targets and about 30 point sources throughout a successful night, with a point-source observation at a similar zenith angle and near in time to any binary observed. In this way, we attempt to minimize both seeing differences as well as residual atmospheric dispersion between the binaries and their calibrators.
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and detector orientation relative to celestial coordinates were measured on each run. To get a precise measure of the pixel scale, we have continued to use a slit mask which mounts to the tertiary mirror baffle support structure. With the mask in place, a bright single star is observed and a normal speckle data file is collected and stored. Using the focal ratio measured by WIYN personnel, the wavelength of the observation, measurements of the separations of the slits on the mask, and the spacing of the line-like features produced in the average spatial frequency power spectrum of the speckle patterns produced, it is possible to determine the pixel scale from first principles. We determine the wavelength of the observation by combining the filter transmission curve, detector quantum efficiency curve versus wavelength, and the spectral type of the star observed. Flux as a function of wavelength for a given spectral type is taken from the Pickles spectral catalogue (Pickles 1998) in order to complete the calculation. In examining the results values obtained in this way for all runs, the stability of the scale values is not high enough to permit averaging over multiple runs, and so we apply each scale value obtained to all the data for the corresponding run. The detector orientation relative to celestial coordinate system is determined by averaging nightly measures of the position of a star when the telescope is offset in various directions by a small amount, typically 10 arcsec. A sequence of 1 s fullframe exposures is taken of a bright star, with the telescope offset between frames. A 200 × 200 subarray containing the star image is extracted from each frame, and the bias and background are estimated from the edges of the array and removed. Finally, the star image is centroided to arrive at the detector coordinates corresponding to a given telescope position. The typical sequence yields six independent measures of the detector orientation plus an estimate of the drift of the star imaged (if any) over the course of the data collection. (Our observations are taken without autoguiding.) The scatter of the six measures then serves as a basic estimate of the uncertainty of the calibration for the night; typical numbers are several tenths of a degree. Excluding the run using the Mini-Mosaic camera, we find in examining the data from all runs that the offset angle was relatively constant throughout 2001, then changed and was stable at a new value for 2002 and 2003, and then changed again but was stable from 2004 through 2006. These three distinct values for the offset angle are explained by the fact that the RYTSI optics were realigned prior to the first run of 2002, and again after the use of the Mini-Mosaic camera in February of 2004. Therefore, we have averaged all the results for each of these time periods to give higher precision on the orientation angles used. Table 1 shows the final scale and offset angles and their estimated uncertainties applied to the data subsequently presented. As noted there, during the 2003 August run, we experimented with the use of the atmospheric dispersion correction (ADC) system in the WIYN port, which was normally stowed during our observations. Having these elements in the beam changes the scale, so that two scale determinations were made during the run, one with the ADC system in place and one without. Also, since we did not have mask observations for the 2005 July and August data, we applied the scale obtained during December 2005 to the observations from these runs. 2.2. Reduction Method
2.1. Pixel Scale and Orientation Except on two occasions where weather did not permit (namely the 2005 July and August runs), the pixel scale
In analyzing data obtained from the RYTSI camera, the first task is to identify the locations of each speckle pattern on the CCD frame. The algorithm used is described in Meyer et al.
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Run 2001 Jun/Jul 2001 Oct 2002 Apr 2002 Oct 2003 Jul 2003 Aug, WIYN ADCs out 2003 Aug, WIYN ADCs in 2004 Feb (Mini-Mo) 2004 Dec 2005 Jul, Aug, Dec 2006 Feb 2006 Jul
Scale (mas pixel−1 ) Orientation angle 27.802 ± 0.056 27.857 ± 0.056 27.108 ± 0.023 27.687 ± 0.085 27.730 ± 0.026 27.610 ± 0.000 26.283 ± 0.040 32.894 ± 0.190 27.844 ± 0.169 28.097 ± 0.016 28.196 ± 0.011 28.367 ± 0.061
176.74 −3.26 2.41 2.41 2.41 2.41 2.41 −88.61 −0.65 −0.65 −0.65 −0.65
± ± ± ± ± ± ± ± ± ± ± ±
0.17 0.18 0.22 0.22 0.22 0.22 0.22 0.15 0.36 0.36 0.36 0.36
(2006), but for completeness we summarize the four essential steps of the process here. First, a copy of the full CCD frame is made and smoothed. This is done by applying a spatial filter to the image that retains image features similar in size to that of the seeing width. The resulting image is thresholded so that pixels with values above the chosen value are assigned a new value of 1 and pixels below the threshold value are assigned a new value of zero. This creates an image that has connected regions of value 1 where speckle patterns exist. Next, the midpoint in both x and y for each connected region is calculated, and this coordinate pair marks the center of the speckle pattern. Finally, in the original image, a box of size 128 × 128 pixels is extracted centered on each coordinate pair, and these are stored in a threedimensional image stack of size 128 × 128 × N where N is the number of speckle patterns found in the original data file. One additional step not included in Meyer et al. is that the read noise is estimated from the oversan region of each frame (or, in the case of Mini-Mosaic data, from a region of the frame containing no signal), and stored for read noise bias subtraction as described below. Once the image stack is created, the subsequent analysis proceeds along the same lines as reported in previous papers, for example Horch et al. (1997). In brief, the average spatial frequency power spectrum is computed for the images, and the read noise and photon bias terms are subtracted. This is done for both binary data and point-source data. Power spectra are also visually inspected at this point and compared with the pointsource calibrator candidate to make sure that the match appears acceptable. If there appears to be any residual bias as judged from the high-frequency corners of the power spectrum (where little or no signal exists), then a bias correction is calculated and applied. Each binary power spectrum is then divided by the power spectrum for its point-source calibrator. The result in theory yields a pure cosine-squared function for a binary star, where the fringe spacing and orientation are related to the binary separation and position angle. The data in the Fourier plane are masked with a roughly annular region, setting all values outside the annulus to zero. The inner radius of the annulus is chosen to exclude the low-frequency contribution of the seeing disk, and the outer edge is chosen to follow a contour of constant signal to noise (thus excluding the high-frequency region of the power spectrum, which is noise dominated). The best-fit function is determined and the astrometric results are obtained. The fringe depth is related to the intensity ratio of the two stars, and so this is also determined by the fitting procedure and converted to a magnitude difference. Two low-order subplanes of the image bi-spectrum are also calculated, so that a reconstructed image of the object can be
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obtained. This permits the determination of the quadrant of the secondary, which is ambiguous in the standard power-spectrum analysis. We have also experimented with a “brightest speckle” shift-and-add algorithm; we found that both methods agreed on the quadrant in 52 out of 53 data files tested. The only case where there was disagreement was a case of a relatively small magnitude difference, 0.32 mag. We chose to use the bi-spectral method, as we have in our previous papers in this series, because it produced cleaner images which allowed for easier identification of faint companions. In Meyer et al. (2006), no flat fielding was done to produce the first results with the RYTSI camera shown there. The difficulty with the RYTSI approach is that the camera has been designed to produce a highly-vignetted field of view on the CCD detector. This is desirable so that light from a wide companion star (not of interest in speckle observations) will be significantly attenuated and will not masquerade in the grid of speckle images as a closer companion. Therefore, when considering how to flat field the speckle data, the optics of the RYTSI box can be expected to produce low-frequency variations in sensitivity, but these will shift in position as the speckle images are directed to different portions of the CCD. On the other hand, high-frequency pixelto-pixel variations may exist which are detector specific and therefore fixed in position. Using a substantial subset of the data presented here (approximately 300 observations), we studied the relative astrometric and photometric accuracy obtained using three different flat-fielding schemes, one of which was simply not to flat field, as in Meyer et al. The results showed that the two other flat-field algorithms did not improve the astrometric or photometric precision of the data, and in fact marginally decreased the accuracy in both cases. This can be explained as follows. Since a large number of speckle images are placed over the active area of the CCD, the pixel-to-pixel sensitivity variations are essentially averaged out. Low-spatialfrequency variations, on the other hand, if larger than the seeing disk, are rendered insignificant by the analysis procedure, which explicitly removes low frequencies prior to fringe fitting. Therefore, any flat-fielding scheme with finite signal-to-noise ratio (S/N) in the flats obtained does little if anything to improve the astrometric and and photometric results. We therefore did not flat field the data in the analysis that produced the results below. 3. RESULTS Table 2 contains our main body of results. The column headings give (1) the Aitken Double Star (ADS), Bright Star Catalogue (i.e., Harvard Revised (HR)), or Durchmusterung (DM) number of the object; (2) the Discoverer Designation; (3) the Draper Catalogue (HD) number; (4) the Hipparcos Catalogue number; (5) the Washington Double Star (WDS) number, which is the same as the right ascension and declination for the object in 2000.0 coordinates; (6) the Besselian date of the observation; (7) the position angle (θ ) of the secondary star relative to the primary, with north through east defining the positive sense of θ ; (8) the separation of the two stars (ρ), in arcsec; (9) the magnitude difference of the pair (∆m); (10) the center wavelength of the filter used for the observation (in nanometers); and (11) the full width at half-maximum (FWHM) of the filter transmission (in nanometers). In some cases, no magnitude difference measure is presented; this is due to the quality cut used for the photometry and described below in Section 3.2. Position angles have not been precessed, and are therefore appropriate for the epoch of the observation shown.
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(a)
(b)
Figure 1. (a) Seeing and seeing difference histograms for the measures in Table 2. (b) Seeing as a function of epoch, showing the variation of seeing within individual runs and from run to run. Table 2 Double Star Speckle Measures HIP
WDS (α, δ J2000.0)
Date (2000+)
θ (◦ )
ρ ( )
∆m
λ ∆λ (nm) (nm)
BD+51 3769 HDS 3 225064 250 ADS 51 HU 1201AB 39 461 461 BD+08 5172 HDS 7 126 510 510 510
00031 + 5228 00055 + 3406 00055 + 3406 00061 + 0943 00061 + 0943 00061 + 0943
2003.6344 2005.5974 2005.5975 2003.6288 2003.6288 2003.6288
335.1 305.6 304.7 213.9 217.0 37.8
0.316 0.236 0.248 0.073 0.073 0.072
2.62 1.70 1.68 0.36 0.21 0.00
541 698 698 550 754 698
HR, ADS DM, etc
Discoverer designation
HD
88 39 39 40a 44a 39b
Notes. a Quadrant ambiguous, but consistent with previous measures in the Fourth Interferometric Catalog. a? Quadrant ambiguous, but possibly consistent with previous measures in the Fourth Interferometric Catalog. b Quadrant ambiguous, but inconsistent with previous measures in the Fourth Interferometric Catalog. b? Quadrant ambiguous, but possibly inconsistent with previous measures in the Fourth Interferometric Catalog. c Quadrant is inconsistent with previous measures in the Fourth Interferometric Catalog. c? Quadrant is possibly inconsistent with previous measures in the Fourth Interferometric Catalog. (This table is available in its entirety in machine-readable and Virtual Observatory (VO) forms in the online journal. A portion is shown here for guidance regarding its form and content.)
Seventeen of the systems listed in Table 2 have no previous position angle and separation measures in the Fourth Interferometric Catalog (Hartkopf et al. 2001b). The majority of these are systems that appear as “suspected” binaries in the Hipparcos Catalogue (ESA, 1997), although occasionally we have found a previously unknown companion to one of our point source calibration objects. We propose discoverer designations of “YSC” (Yale-Southern Connecticut) 1 through 17 for these objects. In the course of preparing Table 2, three minor errors were found in previous papers in this series. First, our 1997.5202 observation of COU 1145 (=HIP 87204 = WDS 17490+3704) does not show the correct separation in Horch et al. (1999). The correct value is 0.116 arcsec, and the error was due to the fact that the observation was not taken with the low-magnification as stated in the paper, but rather at the 2.5× setting. Second, the epochs listed for two observations discussed in Horch et al. (2006) inadvertently had digits reversed in Table 4 of that paper. The values for the observations of MCA 13 (=HIP 19009 = WDS 04044+2406) should read 1997.9672, not 1997.9762.
In Figure 1, we show seeing information for the measures in Table 2. Seeing generally varied between 0.5 and 1.3 arcsec for our observations. Figure 1(a) shows a seeing histogram as well as a histogram for the difference in seeing between each binary and its point-source calibrator. The seeing difference is less than 0.1 arcsec for 78% of the observations in Table 2 (79% for those observations where a magnitude difference is reported), and it is less than 0.2 arcsec for 97% of observations (also 97% for those observations where a magnitude difference is reported). Figure 1(b) shows the seeing measures as a function of epoch. During most runs, a range of 0.5 arcsec or more is seen. In Figure 2, we plot other basic properties of the measures in Table 2. Figure 2(a) shows the magnitude difference obtained as a function of separation. The envelope of this plot may be read as indicating a lower limit to the detectability of faint companions. In Figure 2(b), the magnitude difference is shown as a function of the V magnitude of the system, indicating a limiting magnitude of approximately V = 10 for the front-illuminated CCD used for the vast majority of the observations.
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(a)
(b)
Figure 2. (a) Magnitude difference as a function of separation for the measures listed in Table 2. While a handful of separation measures above 1 arcsec appear in Table 2, the plot has been truncated to clearly show the behavior at sub-arcsecond separations. (b) Magnitude difference as a function of system V magnitude for the measures listed in Table 2. Table 3 Orbits Used for the Measurement Precision Study Discoverer designation FIN 342Aa ANJ 1Aa FIN 331Aa BU 101 FIN 325 STF 1728AB JEF 1 HU 1176AB MCA 47 A 88AB BU 151AB STT 535AB BU 163AB HO 295AB A 417AB MLR 4
HIP
WDS
20661 24608 29850 38382 38474 64241 75695 83838 84949 91394 101769 104858 105200 111805 113996 116849
04256+1556 05167+4600 06171+0957 07518−1354 07528−0526 13100+1732 15278+2906 17080+3556 17217+3958 18384−0312 20375+1436 21145+1000 21186+1134 22388+4419 23052−0742 23411+4613
Coordinate(s) used Grade ρ ρ ρ θ θ θ ρ ρ θ, ρ ρ θ, ρ θ, ρ θ θ ρ ρ
1 1 1 2 2 1 2 1 2 1 1 1 2 2 1 2
Reference McAlister et al. (1988) Bagnuolo and Hartkopf (1989) Hartkopf et al. (1996) W. I. Hartkopf (2001, private communication) Hartkopf et al. (1996) Hartkopf et al. (1989) McAlister et al. (1992) Hartkopf et al. (1989) Scarfe et al. (1994) Hartkopf et al. (1989) W. I. Hartkopf (2001, private communication) Hartkopf et al. (1996) Fekel et al. (1997) Hartkopf et al. (1996) Hartkopf et al. (1996) Hartkopf et al. (1996)
Table 4 Astrometric Measurement Precision Object type All Speckle orbits with δρ < 1.5 mas Grade 1 orbits with δρ < 1.5 mas Multiple measures All Speckle orbits with δθ < 0.7◦ Speckle orbits with δθ < 0.7◦ Multiple measures
Parameter
Average residual
ρ ρ ρ θ θ θ
−0.11 ± 0.40 mas −0.04 ± 0.48 mas ... −0.09 ± 0.10◦ +0.03 ± 0.09◦ ...
3.1. Astrometric Accuracy and Precision In order to characterize the astrometric precision, a subset of objects with extremely well-known orbits were chosen from the Sixth Catalog of Visual Orbits of Binary Stars (Hartkopf et al. 2001a). The orbits selected are shown in Table 3. The basic criteria for inclusion in Table 3 are the existence of a large body of speckle observations and a recent orbit calculation where most of these data have been included. The majority of such objects have Grade 1 (definitive) orbits in the Sixth Catalog, though we also included several other objects that have a large body of recent speckle data for a further high quality
rms deviation from Number of average residual measures 2.81 ± 0.28 mas 2.89 ± 0.34 mas 2.42 ± 0.30 mas 0.88 ± 0.07◦ 0.61 ± 0.06◦ 0.48 ± 0.07◦
49 36 83 78 48 83
comparison. In all cases, the authors who calculated the orbits stated uncertainties in the orbital elements; these are used to estimate uncertainties in the ephemeris positions using standard error propagation methods. Using measures from Table 2, we form residuals in position angle and separation (observed value minus the ephemeris prediction), and we plot these results in Figure 3. The error bars shown in these figures are the estimated uncertainties in the ephemeris positions. Table 4 shows the average residuals and root mean square (rms) deviation from the average residual for these cases. Lines 1, 2, 4, and 5 specifically refer to the speckle orbit study. There are no significant deviations in position
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(a)
(b)
Figure 3. (a) Residuals in position angle when comparing measures in Table 2 with the predicted position angle in the case of high-quality orbits. The open circles are orbits with predicted uncertainties
