Limits on Unresolved Planetary Companions to White Dwarf Remnants of 14 Intermediate-Mass Stars
aa r X i v : . [ a s t r o - ph . S R ] S e p Limits on Unresolved Planetary Companions to White DwarfRemnants of 14 Intermediate-Mass Stars
Mukremin Kilic , , Andrew Gould , and Detlev Koester ABSTRACT
We present
Spitzer
IRAC photometry of white dwarf remnants of 14 stars with M = 3 − M ⊙ . We do not detect mid-infrared excess around any of our targets.By demanding a 3 σ photometric excess at 4.5 µ m for unresolved companions, werule out planetary mass companions down to 5, 7, or 10 M J for 13 of our targetsbased on the Burrows et al. (2003) substellar cooling models. Combined withprevious IRAC observations of white dwarf remnants of intermediate-mass stars,we rule out ≥ M J companions around 40 white dwarfs and ≥ M J companionsaround 10 white dwarfs. Subject headings: infrared: stars — planetary systems — stars: low-mass, browndwarfs — white dwarfs
1. INTRODUCTION
Radial velocity, transit, and microlensing searches are succesful in finding planets aroundstars less massive than 2 M ⊙ (see Udry & Santos 2007; Marcy et al. 2005; Gould 2009; Mazeh2009, and references therein). However, these techniques have limitations for higher massstars. The main problem is that intermediate- and high-mass stars are big, they are rare,and they do not have many usable lines for radial velocity studies. The short ( ≤ Smithsonian Astrophysical Observatory, 60 Garden Street, Cambridge, MA 02138, USA Department of Astronomy, Ohio State University, 140 W. 18th Ave., Columbus, OH 43210, USA Institut f¨ur Theoretische Physik und Astrophysik, University of Kiel, 24098 Kiel, Germany Spitzer Fellow; [email protected] . M ⊙ ) WDs, increasing the contrast for substellar companionseven further. If planetary systems survive the late stages of stellar evolution, massive planetsshould be detectable around WDs (Burleigh et al. 2002).Detailed dynamical simulations of the evolution of the Solar system by Duncan & Lissauer(1998) show that the giant planets are likely to remain in stable orbits for more than 10 Gyrafter the Sun becomes a WD. The discovery of an M sin i = 3 . M J planet around the sdB starV391 Pegasi by Silvotti et al. (2007) and a candidate ≥ . M J planet around the WD GD66 by Mullally et al. (2009) indicate that at least some planets survive post main-sequenceevolution.There are now more than a dozen WDs known to host circumstellar debris disks (seeFarihi et al. 2009; Jura et al. 2009; Kilic et al. 2008, and references therein). These disksare most likely formed by tidally disrupted asteroids (Jura 2003), and at least in one case(GD 362), the composition of the accreted matter most resembles that of the Earth/Moonsystem (Zuckerman et al. 2007). The atmospheric composition of other disk-polluted WDshas not yet been analyzed in such detail. Farihi et al. (2009) estimate that at least 1% to3% of WDs with cooling ages less than about 0.5 Gyr harbor debris disks. There is indirectevidence from these debris disks that at least a few percent of WDs have remnant planetarysystems (Jura 2003, 2006).Laws et al. (2003) and Fischer & Valenti (2005) find a significant trend in the frequencyof planets with increasing stellar mass. The frequency of Jovian-planets ( M > . M J and a < . . − . M ⊙ stars (Johnson et al.2007). Of course, the analogs of all of these planets around higher mass stars would mostlikely be swallowed in the asymptotic giant branch (AGB) phase (Villaver & Livio 2007),and the frequency of planets around massive stars at wider separations is poorly constrainedat present. In the planet formation models of Kennedy & Kenyon (2008), the fraction ofstars with giant planets shows a steady increase with mass up to 3 M ⊙ . In addition, the massof the planets and the width of the regions where they form are predicted to increase withstellar mass. Observations of massive WDs can test these models.Gould & Kilic (2008) identified 49 young ( ≤ Spitzer
IRAC observations of 14 stars from that sample. Our observations arediscussed in Section 2, while the spectral energy distributions and the limits on planetarycompanions are discussed in Section 3. 3 –
2. OBSERVATIONS
We selected our targets from the DA WDs found in the Palomar-Green Survey (Liebert et al.2005). The advantage of this sample selection is that Liebert et al. (2005) provide T eff , log g ,mass, cooling age, and distance estimates for all these stars. We selected the 14 bright-est WDs with M W D = 0 . − . M ⊙ from that sample. The physical parameters of ourtargets, including the initial-mass estimates derived using the initial-final mass relation ofKalirai et al. (2008), are presented in Table 1. The initial-masses of our targets range from2.8 to 4.6 M ⊙ . The use of different initial-final mass relations results in a ≈
10% difference inthe initial-mass estimates. We estimate the main-sequence (MS) lifetimes using the equation t MS = 10( M MS M ⊙ ) − . Gyr (Wood 1992). The total (WD + MS) ages of our targets are alsopresented in Table 1. Four targets are about 300 Myr old, and the remaining targets haveages ≈ Spitzer
FellowshipProgram 474. We obtained 3.6, 4.5, 5.8, and 7.9 µ m images with integration times of 30or 100 seconds per dither, with five or nine dithers per object. We use the IRAF PHOTand IDL astrolib packages to perform aperture photometry on the individual BCD framesfrom the latest available IRAC pipeline reduction. Since our targets are relatively faint, weuse the smallest aperture (two pixels) for which there are published aperture corrections.Following the IRAC calibration procedure, corrections for the location of the source in thearray are taken into account before averaging the fluxes of each of the dithered framesat each wavelength. Channel 1 (3.6 µ m) photometry is also corrected for the pixel-phase-dependence. The results from IRAF and IDL reductions are consistent within the errors.The photometric error bars are estimated from the observed scatter in the five (or nine)images corresponding to the dither positions. We also add the 3% absolute calibration errorin quadrature. Finally, we divide the estimated fluxes by the color corrections for a Rayleigh-Jeans spectrum (Reach et al. 2005). These corrections are 1.0111, 1.0121, 1.0155, and 1.0337for the 3.6, 4.5, 5.8, and 7.9 µ m bands, respectively.We present the IRAC photometry of our targets in Table 2. All of our targets aredetected in 2MASS, at least in the J and H bands. All but two of our targets also haveSloan Digital Sky Survey (SDSS) photometry available. PG 0852+659 and PG 1335+701are not covered in the SDSS Data Release 7 area. We obtained V and I band photometryof PG 1335+701 using the MDM 2.4m telescope equipped with the Echelle CCD. Two setsof 120 s exposures were obtained in each filter on UT 2009 April 5. We use observationsof the standard star field PG 1323 −
086 (Landolt 1992) to calibrate the photometry. Thephotometric reductions were performed by J. Thorstensen, and kindly made available tous. PG 1335+701 has V = 15.29 mag and I = 15.54 mag. We estimate an internal accu- 4 –racy of 0.01 mag for the photometry, but to be conservative, we adopt errors of 0.03 mag.Liebert et al. (2005) provide an estimate on the V magnitude of PG 0852+659 using pho-tographic B − band photometry. However, the error in the photographic magnitude is largerthan 0.3 mag, and therefore we do not use it in our analysis. PG 0852+659 is the only WDin our sample without accurate optical photometry.
3. RESULTS
Figure 1 presents the spectral energy distributions of four WDs with ages ≈
300 Myr.This youth is important; young planets will be brighter than their older counterparts andtherefore it should be easier to find planets around relatively young WDs. The SDSS pho-tometric zero points differ slightly from the AB convention (Eisenstein et al. 2006). We usethe corrections given in Eisenstein et al. (2006) to convert the photometry to the AB system.We calculate synthetic spectra for our targets using the model atmosphere code describedby Koester (2009) and the best-fit T eff and log g values from Liebert et al. (2005). We per-form synthetic photometry on these WD models using the appropriate transmission curvesfor the SDSS, 2MASS, and IRAC filters. The resulting fluxes are then compared with theobservations to find the normalization factor for the WD models. We weight fluxes by theirassociated error bars. The solid lines in Figure 1 present the appropriate WD model foreach star, which is normalized to match the observations. Optical and near-infrared pho-tometry helps us to constrain the predicted mid-infrared photospheric fluxes for WDs. The4.5 µ m photometry of all four WDs presented in this figure is consistent with the predictedphotospheric flux from WDs within 1 σ ; none of them show excess mid-infrared flux.We use the synthetic spectra for 300 Myr old 1-25 M J planets (Burrows et al. 2003) toput upper limits on companions that would escape detection. The overall agreement betweenthe Hubeny & Burrows (2007) models and the spectral energy distributions of cold browndwarfs in the T dwarf range suggests that the planet models that we use are appropriate forthe colder planets expected around WDs. The models by A. Burrows agree reasonably wellwith the observed hot Jupiter emission spectra in the Spitzer
IRAC bands (Burrows et al.2008, H. Knutson 2009, priv. comm.). However, the predicted fluxes can vary by a factor oftwo due to temperature inversions in the atmosphere and water being observed in emissionor absorption. These temperature inversions would make the infrared fluxes higher, and theywould make it easier to find planets.The dotted lines in Figure 1 show the combined flux from each WD plus 10, 7, 5, 2, and1 M J companions (from top to bottom, Burrows et al. 2003), respectively. The red crossmarks the 3 σ upper limit of the 4.5 µ m photometry. By demanding a 3 σ excess from any 5 –possible companion, we exclude ≥ M J planets around PG 1038+634, PG 1051+274, andPG 1446+286. We exclude ≥ M J companions around PG 1335+701. Of course, thesemass limits are model-dependent.Figure 2 presents the spectral energy distributions of the remaining 10 WDs in oursample. All of these WDs are older than about 600 Myr. We use 1 Gyr old planet modelsto constrain the contribution from possible companions. As in Figure 1, the solid lines showthe best-fit WD models. The dotted lines show the combined flux from each WD plus 1 Gyrold 25, 20, 15, 10, 7, 5, 2, and 1 M J companions (from top to bottom), respectively. None ofthe WDs in Figure 2 show significant flux excesses in the mid-infrared. We rule out planetsmore massive than 5, 7, or 10 M J around these targets with 3 σ confidence.Table 3 presents the search radius and the unresolved companion limits for each star.Our two pixel search radius in IRAC images corresponds to 50 AU for the nearest, and 260AU for the most distant WD in our sample. We expect the mass loss process to enlarge theplanetary orbits by a factor of M MS /M WD (Zuckerman & Becklin 1987). We use the initial-final mass relation of Kalirai et al. (2008) to estimate the ratio of MS and WD masses, i.e.,the orbital expansion factor. The above search radii allow us to probe the WDs’ progenitorsfor planets to 11 AU for the nearest WD and to 55 AU for the most distant one. For theseparations probed (see Table 3) at 13 of our targets, we do not find any planets more massivethan 10 M J , according to Burrows et al. (2003) models. For five of our targets, there are noplanets more massive than 5 M J within the probed inner regions.We extend our search to partially-resolved companions by increasing the size of thephotometric aperture used in our analysis. Also, it may be possible to exclude resolvedcompanions at larger radii by the lack of 4.5 µ m point sources around our targets. Oneof our targets, PG 1307+354, has two nearby sources that contaminate the photometry inapertures larger than two pixels. For the remaining 13 targets, the differences between threepixel aperture photometry and two pixel aperture photometry are relatively small ( ≤ − M J ) to be detected around some of our targets. Using athree pixel aperture and accounting for orbital expansion, we would have detected a massiveplanet at 24 AU around 13 of our targets, a planet at 38 AU around six of our targets, anda planet at 68 AU around only three of our targets. 6 –
4. DISCUSSION
None of the stars in our sample show mid-infrared flux excess from brown dwarfs orplanetary mass companions. Our data rule out ≥ M J companions for five WDs and ≥ M J companions for 13 WDs.There have been many other searches for substellar and planetary mass companions toWDs. In fact, the first candidate brown dwarf was found around a WD more than 20 yearsago (Becklin & Zuckerman 1988). However, only a few more WD + brown dwarf systemshave been discovered since then (Farihi & Christopher 2004; Maxted et al. 2006; Steele et al.2009). Farihi et al. (2005) find that less than 0.5% of WDs have brown dwarf companions.Hogan et al. (2009) performed a J − band proper motion survey of 23 WDs with Gemini,and found that ≤
5% of WDs have substellar companions. In addition, the near- and mid-infrared searches by Debes et al. (2005, 2007) and Friedrich et al. (2006) did not reveal anysubstellar companions to WDs.
Spitzer
IRAC currently provides the best opportunity todetect planetary mass companions around WDs. An IRAC survey of 124 nearby WDs byMullally et al. (2007) did not find any planets. However, accurate mass and age estimatesare not available for the majority of the stars in their sample, and a detailed analysis wouldbe required to put reliable limits on possibly hidden companions.Farihi et al. (2008) present IRAC observations of 48 WDs including 31 WDs youngerthan 1 Gyr (MS + WD cooling age). They use blackbody models to predict the 4.5 µ mphotospheric flux from their targets and search for 15% excess flux at 4.5 µ m. Of course,the use of WD-atmosphere models (as in our study) would be more accurate, but blackbodymodels are sufficient for finding 15% excesses around relatively hot WDs. They rule out ≥ M J companions around 27 of their targets. The addition of 13 stars presented in thispaper brings the total sample size to 40. None of the 40 stars in the combined sample showinfrared excess due to substellar companions more massive than 10 M J . For a limit of 5 M J ,there are a total of 10 WDs in both studies. While no planets are detected (i.e., f = 0), itremains conceivable that the frequency is non-zero.Due to small sample size, we use a binomial probability distribution to derive statis-tical uncertainties. The probability, P ( f ), that a survey of N stars will detect n com-panions, when the true frequency of companions is f is given by (Burgasser et al. 2003;McCarthy & Zuckerman 2004): P n ( f ) = f n (1 − f ) N − n N !( N − n )! n ! . (1) 7 –For N = 40 and n = 0, the probability distribution peaks at zero. Since the distributionis not symmetric about its maximum value, we report the range in frequency that delimits34% and 68% of the integrated probability function as the mean frequency and error bars,respectively. These error bars are equivalent to 1 σ limits for a Gaussian distribution. Wefind that the frequency of ≥ M J companions to WDs is 1 . +1 . − . %. This is consistent withzero.Figure 3 displays planet versus host star mass for all known extrasolar planets detectedby the radial velocity, astrometric, transit, microlensing, and direct imaging searches as of2009 April (based on the Extrasolar Planets Encyclopedia ). The parameter space that isruled out by IRAC observations of the WD remnants of intermediate-mass stars presented inthis work and Farihi et al. (2008) is also shown. We note that this figure excludes seven starswith M W D > . M ⊙ from the Farihi et al. (2008) study. The initial-final mass relation ofKalirai et al. (2008) implies that these WDs are the descendants of stars with M ≥ . M ⊙ .However, the initial-final mass relation is steep near the top end of the WD mass function(Williams et al. 2009), and the errors are relatively large. The initial-final mass relation ofWilliams et al. (2009) implies an initial-mass of ≈ M ⊙ for a 1 . M ⊙ WD. Even though theseseven stars are not shown in Figure 3, they are included in our statistical analysis of 40 starspresented above.About 60% of the Sun-like (0 . − . M ⊙ ) stars with RV detected planets have planet/hoststar mass ratios ≥ M J M ⊙ (based on the Extrasolar Planets Encyclopedia). This fraction goesdown to 33% for planet/host star mass ratios greater than 2 M J M ⊙ . Johnson et al. (2007) findthat the Jovian-planet frequency is 4.2% for Sun-like stars and it is 8.9% for 1 . − . M ⊙ stars.This trend is expected to continue for higher mass stars up to 3 − M ⊙ . Kennedy & Kenyon(2008) predict that about 20% of 3 − M ⊙ stars have at least one gas giant planet. Theyalso find that the isolation masses are larger for more massive stars, and the giant planetsthat eventually form should also be bigger (see Fig. 2 and 3 in Kennedy & Kenyon 2008).Assuming that the planet/host star mass ratios are similar for Sun-like and intermediatemass stars, we expect 6.6% of 3 − M ⊙ stars to have planets with planet/host star massratios greater than 2 M J M ⊙ . We probe this regime for 6 stars in our sample. Farihi et al. (2008)reach the same limit for 20 stars, including 10 stars with M W D > . M ⊙ . All 26 stars have M MS ≥ M ⊙ . The probability of a null detection among a sample of 26 stars when the(predicted) frequency is 6.6% is 17% (see equation 1). Thus, the null detections reported inthis paper and Farihi et al. (2008) are suggestive, but do not directly conflict with previousdetections of planets around intermediate-mass stars. Our sample is not large enough to put http://exoplanet.eu/ M ⊙ , irradiation overcomes accretion as the stars reach the mainsequence relatively quickly. This pushes the snow line to 10 −
15 AU and makes formationof cores difficult (Ida & Lin 2005; Kennedy & Kenyon 2008).2 - The planets may not survive the red giant and the AGB phases. Villaver & Livio(2007) find that planets within 5.3 AU of 5 M ⊙ stars will be engulfed during the AGB phase.Of course, avoiding engulfment may not be enough for a planet to survive because a corpus-cular drag resulting from its interaction with the stellar wind would decrease the planetaryorbits. However, Duncan & Lissauer (1998) show that the planets will only move inwardslightly due to the corpuscular drag. Villaver & Livio (2007) suggest that the maximumstellar radius is reached only for a brief period of time, and the planetary orbits would havebeen expanded due to mass loss by that time. Hence, they find that the tidal drag forcesare negligible for orbital radii larger than the maximum stellar radius reached in the AGBphase. Since the search radii for planets around the progenitors of our sample of WDs rangesfrom 11 to 55 AU, the engulfment of the planets during the AGB phase cannot explain theapparent lack of high mass planets around our targets.3 - The orbits of the planets around massive stars may become unstable during the latestages of stellar evolution. Debes & Sigurdsson (2002) suggest that planets around WDsmay become unstable to close approaches with each other and the entire system may becomedynamically young. Duncan & Lissauer (1998) find that the crossing times for planetaryorbits depend on ( M WD /M MS ) . . Even though this timescale is relatively long for the Sun( ∼
10 Gyr), the larger mass-reduction factor for more massive stars means that the timescalefor unstable close approaches is less than 1 Gyr for planets around M ≥ . M ⊙ stars. Inaddition, since orbital semi-major axes grow during the mass loss phase, some planets willpass through resonances that could create or enhance instabilities in the system. However,detailed simulations including the effects of orbital expansion on the stability of planetsaround massive WDs are not currently available. 9 –
5. CONCLUSIONS
There is only one known WD that has a possible planetary companion, GD 66. Thesignal from an M sin i = 2 . M J planetary companion in a 4.5 yr orbit is detected in timingmeasurements of this pulsating WD star (Mullally et al. 2009). However, a complete orbithas not been observed yet, and the detection remains provisional. Spitzer
IRAC observationsof GD 66 did not reveal any significant mid-infrared flux excess. Based on the substellar cool-ing models by Burrows et al. (2003) and looser detection criteria compared to Farihi et al.(2008) and our study, Mullally et al. (2009) place an upper limit of 5 − M J on the mass ofthe companion.So far, no other search, including the IRAC surveys by Mullally et al. (2007), Farihi et al.(2008), and this paper has been succesful in finding planetary companions to WDs. Theproper motion surveys in the near-infrared have ruled out resolved companions to two dozenWDs (Hogan et al. 2009). However, a proper motion survey of the majority of the WDs ob-served with IRAC has not been performed yet. Such a survey will be valuable for searchingfor resolved companions to WDs. Future studies with ground based telescopes using Adap-tive Optics instruments and with the James Webb Space Telescope (JWST) may provide thefirst answers to whether Jupiter mass planets survive around WDs or not.Support for this work was provided by NASA through the Spitzer Space TelescopeFellowship Program, under an award from Caltech. A.G. was supported by NSF grant AST-0757888. We thank the referee, J. Farihi, for a detailed and constructive report. We alsothank S. Kenyon for a careful reading of this manuscript.
REFERENCES
Becklin, E. E., & Zuckerman, B. 1988, Nature, 336, 656Burgasser, A. J., Kirkpatrick, J. D., Reid, I. N., Brown, M. E., Miskey, C. L., & Gizis, J. E.2003, ApJ, 586, 512Burleigh, M. R., Clarke, F. J., & Hodgkin, S. T. 2002, MNRAS, 331, L41Burrows, A., Sudarsky, D., & Lunine, J. I. 2003, ApJ, 596, 587Burrows, A., Budaj, J., & Hubeny, I. 2008, ApJ, 678, 1436Debes, J. H. & Sigurdsson, S. 2002, ApJ, 572, 556 10 –Debes, J. H., Sigurdsson, S., & Woodgate, B. E. 2005, ApJ, 633, 1168Debes, J. H., Sigurdsson, S., & Hansen, B. 2007, AJ, 134, 1662Duncan, M. J., & Lissauer, J. J. 1998, Icarus, 134, 303Eisenstein, D. J., et al. 2006, ApJS, 167, 40Farihi, J., & Christopher, M. 2004, AJ, 128, 1868Farihi, J., Becklin, E. E., & Zuckerman, B. 2005, ApJS, 161, 394Farihi, J., Becklin, E. E., & Zuckerman, B. 2008, ApJ, 681, 1470Farihi, J., Jura, M., & Zuckerman, B. 2009, ApJ, 694, 805Fischer, D.A. & Valenti, J. 2005, ApJ, 622, 1102Friedrich, S., Zinnecker, H., Correia, S., Brandner, W., Burleigh, M., & McCaughrean, M.2006, ASP Conference Series, 999Gould, A., & Kilic, M. 2008, ApJ, 673, L75Gould, A. 2009, Astronomical Society of the Pacific Conference Series, 403, 86Hogan, E., Burleigh, M. R., & Clarke, F. J. 2009, MNRAS, in pressHubeny, I., & Burrows, A. 2007, ApJ, 669, 1248Ida, S., & Lin, D. N. C. 2005, ApJ, 626, 1045Ignace, R. 2001, PASP, 113, 1227Johnson, J.A., Butler, R.P., Marcy, G.W., Fischer, D.A., Vogt, S.S., Wright, J.T., & Peek,K.M.G., 2007, ApJ, 670, 833Jura, M. 2003, ApJ, 584, L91Jura, M. 2006, ApJ, 653, 613Jura, M., Farihi, J., & Zuckerman, B. 2009, AJ, 137, 3191Kalas, P., et al. 2008, Science, 322, 1345Kalirai, J. S., Hansen, B. M. S., Kelson, D. D., Reitzel, D. B., Rich, R. M., & Richer, H. B.2008, ApJ, 676, 594 11 –Kennedy, G. M., & Kenyon, S. J. 2008, ApJ, 673, 502Kilic, M., Farihi, J., Nitta, A., & Leggett, S. K. 2008, AJ, 136, 111Koester, D. 2009, Mem. Soc. Astron. Italiana, in pressKoester, D., Rollenhagen, K., Napiwotzki, R., Voss, B., Christlieb, N., Homeier, D., &Reimers, D. 2005, A&A, 432, 1025Landolt, A. U. 1992, AJ, 104, 340Laws, C., Gonzalez, G., Walker, K.M., Tyagi, S., Dodsworth, J., Snider, K., & Suntzeff,N.B. 2003, AJ, 125, 2664L´epine, S., & Shara, M. M. 2005, AJ, 129, 1483Liebert, J., Bergeron, P., & Holberg, J. B. 2005, ApJS, 156, 47Livio, M., Pringle, J. E., & Wood, K. 2005, ApJ, 632, L37Marcy, G., Butler, R. P., Fischer, D., Vogt, S., Wright, J. T., Tinney, C. G., & Jones,H. R. A. 2005, Progress of Theoretical Physics Supplement, 158, 24Marois, C., Macintosh, B., Barman, T., Zuckerman, B., Song, I., Patience, J., Lafreni`ere,D., & Doyon, R. 2008, Science, 322, 1348Maxted, P. F. L., Napiwotzki, R., Dobbie, P. D., & Burleigh, M. R. 2006, Nature, 442, 543Mazeh, T. 2009, IAU Symposium, 253, 11McCarthy, C., & Zuckerman, B. 2004, AJ, 127, 2871Mullally, F., Kilic, M., Reach, W. T., Kuchner, M. J., von Hippel, T., Burrows, A., &Winget, D. E. 2007, ApJS, 171, 206Mullally, F., Reach, W. T., Degennaro, S., & Burrows, A. 2009, ApJ, 694, 327Reach, W. T., et al. 2005, PASP, 117, 978Silvotti, R., et al. 2007, Nature, 449, 189Steele, P. R., Burleigh, M. R., Farihi, J., G¨ansicke, B. T., Jameson, R. F., Dobbie, P. D., &Barstow, M. A. 2009, A&A, 500, 1207Udry, S., & Santos, N. C. 2007, ARA&A, 45, 397 12 –Villaver, E., & Livio, M. 2007, ApJ, 661, 1192Williams, K. A., Bolte, M., & Koester, D. 2009, ApJ, 693, 355Wood, M. A. 1992, ApJ, 386, 539Zuckerman, B., & Becklin, E. E. 1987, ApJ, 319, L99Zuckerman, B., Koester, D., Reid, I. N., & H¨unsch, M. 2003, ApJ, 596, 477Zuckerman, B., Koester, D., Melis, C., Hansen, B. M., & Jura, M. 2007, ApJ, 671, 872
This preprint was prepared with the AAS L A TEX macros v5.2.
13 –Table 1. High Mass WD Targets
Object T eff log g M WD M MS d τ WD τ WD+MS (K) (cm s − ) ( M ⊙ ) ( M ⊙ ) (pc) (Myr) (Myr)PG 0852+659 19070 8.13 0.70 2.8 92 140 900PG 1034+492 20650 8.17 0.73 3.1 80 120 720PG 1038+634 24450 8.38 0.87 4.4 68 100 350PG 1051+274 23100 8.37 0.86 4.3 41 110 380PG 1108+476 12400 8.31 0.80 3.7 46 600 980PG 1129+156 16890 8.19 0.73 3.1 36 220 830PG 1201 −
001 19770 8.26 0.78 3.5 63 160 580PG 1307+354 11180 8.15 0.70 2.8 45 630 1390PG 1310+583 10560 8.32 0.80 3.7 21 910 1290PG 1319+466 13880 8.19 0.73 3.1 37 400 1000PG 1335+701 30140 8.25 0.79 3.6 108 30 420PG 1446+286 22890 8.42 0.89 4.6 47 130 350PG 1515+669 10320 8.40 0.86 4.3 33 1120 1390PG 1550+183 14260 8.25 0.77 3.5 41 390 840
Table 2. IRAC Photometry of High Mass WDs
Object 3.6 µ m 4.5 µ m 5.8 µ m 8.0 µ m Reduction( µ Jy) ( µ Jy) ( µ Jy) ( µ Jy) PipelinePG 0852+659 97.7 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± −
001 123.1 ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
14 –Table 3. Limits on Unresolved Companions
Object Search Radius † Initial Separation † M MS Ruled out †† (AU) (AU) ( M ⊙ ) CompanionsPG 0852+659 221 55 2.8 > M J PG 1034+492 192 46 3.1 > M J PG 1038+634 163 32 4.4 ≥ M J PG 1051+274 98 20 4.3 ≥ M J PG 1108+476 110 24 3.7 ≥ M J PG 1129+156 86 20 3.1 > M J PG 1201 −
001 151 33 3.5 ≥ M J PG 1307+354 108 27 2.8 > M J PG 1310+583 50 11 3.7 ≥ M J PG 1319+466 89 21 3.1 ≥ M J PG 1335+701 259 56 3.6 ≥ M J PG 1446+286 113 22 4.6 ≥ M J PG 1515+669 79 16 4.3 ≥ M J PG 1550+183 98 22 3.5 ≥ M J † For a two pixel aperture. †† These limits are calculated based on the models by Burrows et al. (2003).See the discussion in Section 3.
15 –
Fig. 1.— Spectral energy distributions of four 300 Myr old (total age) WDs. The expectedphotospheric flux from the WDs are shown as solid lines (Koester 2009). The dotted linesshow the expected flux from planetary companions with M = 10 , , , M J (from topto bottom), respectively. The red cross marks the 3 σ upper limit of the 4.5 µ m photometry. 16 – Fig. 2.— Spectral energy distributions of ≈ M = 25 , , , , , , M J (from top to bottom), respectively. The red cross marks the 3 σ upper limit of the 4.5 µ mphotometry. 17 – Fig. 2.— contd. 18 –
Fig. 2.— contd. 19 –Fig. 3.— Planet versus host mass for all known extrasolar planets detected by the RV,astrometry, transits, microlensing, and direct imaging observations as of 2009 April. Thedashed line marks the upper limit for planetary mass objects ( M = 13 M J ). The parameterspace that is ruled out by Spitzer observations of WD remnants of intermediate-mass starsby Farihi et al. (2008) and this work is shown as black and blue lines, respectively. We use3 σσ