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8

XO-3b: A Massive Planet in an Eccentric Orbit Transiting an

F5V Star

Christopher M. Johns–Krull1,2, Peter R. McCullough3, Christopher J. Burke3,

Jeff A. Valenti3, K. A. Janes4, J. N. Heasley5, L. Prato6, R. Bissinger7, M. Fleenor8,

C. N. Foote9, E. Garcia–Melendo10, B. L. Gary11, P. J. Howell4, F. Mallia12, G. Masi13,

T. Vanmunster14

cmj@rice.edu

ABSTRACT

We report the discovery of a massive (Mpsini= 13.02± 0.64 MJ; total mass 13.25± 0.64 MJ),large (1.95 ± 0.16 RJ) planet in a transiting, eccentric orbit (e = 0.260 ± 0.017) around a 10th

magnitude F5V star in the constellation Camelopardalis. We designate the planet XO-3b, and thestar XO-3, also known as GSC 03727-01064. The orbital period of XO-3b is 3.1915426± 0.00014days. XO-3 lacks a trigonometric distance; we estimate its distance to be 260±23 pc. The radiusof XO-3 is 2.13±0.21 R⊙, its mass is 1.41±0.08 M⊙, its vsini = 18.54 ± 0.17 km s−1, and itsmetallicity is [Fe/H] = −0.177 ± 0.027. This system is unusual for a number of reasons. XO-3bis one of the most massive planets discovered around any star for which the orbital period isless than 10 days. The mass is near the deuterium burning limit of 13 MJ, which is a proposedboundary between planets and brown dwarfs. Although Burrows et al. (2001) propose thatformation in a disk or formation in the interstellar medium in a manner similar to stars is amore logical way to differentiate planets and brown dwarfs, our current observations are notadequate to address this distinction. XO-3b is also unusual in that its eccentricity is large givenits relatively short orbital period. Both the planetary radius and the inclination are functionsof the spectroscopically determined stellar radius. Analysis of the transit light curve of XO-3bsuggests that the spectroscopically derived parameters may be over estimated. Though relativelynoisy, the light curves favor a smaller radius in order to better match the steepness of the ingressand egress. The light curve fits imply a planetary radius of 1.25±0.15 RJ, which would correspondto a mass of 12.03 ± 0.46 MJ. A precise trigonometric parallax measurement or a very accuratelight curve is needed to resolve the uncertainty in the planetary mass and radius.

Subject headings: binaries: eclipsing – planetary systems – stars: individual (GSC 03727-01064) – tech-niques: photometric – techniques: radial velocities

1Dept. of Physics and Astronomy, Rice University, 6100Main Street, MS-108, Houston, TX 77005

2Visiting Astronomer, McDonald Observatory, which isoperated by the University of Texas at Austin.

3Space Telescope Science Institute, 3700 San MartinDr., Baltimore MD 21218

4Boston University, Astronomy Dept., 725 Common-wealth Ave., Boston, MA 02215

5University of Hawaii, Inst. for Astronomy, 2680 Wood-lawn Dr., Honolulu, HI 96822

6Lowell Observatory, 1400 West Mars Hill Road,

Flagstaff, AZ, 860017Racoon Run Observatory, 1142 Mataro Court,

Pleasanton, CA 945668Volunteer Observatory, 10305 Mantooth Lane,

Knoxville, TN 379329Vermillion Cliffs Observatory, 4175 E. Red Cliffs Drive,

Kanab, UT 8474110Esteve Duran Observatory, El Montanya, Seva, 08553

Seva, Barcelona, Spain11Hereford Arizona Observatory, 5320 E. Calle Manzana,

Hereford, AZ 8561512Campo Catino Astronomical Observatory, P.O. BOX

1

1. Introduction

There are now over 200 extrasolar planetsknown (http://exoplanets.eu/), and most of thesehave been discovered using the radial velocitytechnique. As a result, most of these systemsyield the minimum mass of the planet (Mpsini),its orbital semi-major axis, and properties of thecentral star. The data to date have dramaticallychanged our appreciation of the diversity of plane-tary systems that can form, and have particularlyfocussed the community’s attention on the rolethat planetary migration plays in the planet for-mation process (see review by Papaloizou et al.2007). As the precision and duration of radialvelocity surveys increases, lower mass planets andplanets in longer orbits continue to be found. Asa result, it has been suggested that the fraction ofstars hosting planets may be as high as 50% (seeUdry et al. 2007). Therefore, it appears that theplanet formation process is relatively efficient.

The generally accepted model of giant planetformation is that of core nucleated accretion (e.g.,Pollack et al. 1996; Bodenheimer et al. 2000;Hubickyj et al. 2005) in which gas giant plan-ets first form a several (∼ 10) earth mass coreof solids via non-elastic collisions in a disk fol-lowed by the runaway accretion of gas from thedisk once the mass of the core is sufficient to ex-ert a strong gravitational influence. Such a modelnaturally predicts a positive correlation betweenthe metallicity of the disk (and parent star) andthe ease with which planets can form: there aremore building blocks to form the solid core inhigher metallicity cases. The expected correlationhas been suspected for some time (e.g., Gonza-lez 1997; Santos et al. 2003), and has now beendemonstrated in an unbiased way (e.g., Santos etal. 2004; Fischer & Valenti 2005) and has beentaken as evidence in support of the core accretionmodel of planet formation. The competing model,gravitational instability in a massive circumstellardisk (e.g., Boss 1997, 2000; Mayer et al. 2002)does not predict such a relationship. Livio andPringle (2003) showed that lower disk metallicity

03016, Guarcino (FR) Italy13Bellatrix Observatory, Via Madonna de Loco, 47 03023

Ceccano (FR) Italy14CBA Belgium Observatory, Walhostraat 1A, B-3401

Landen, Belgium

reduces the efficiency of Type II migration whichmay account for some of the increased probabilityfor finding a planet around more metal rich stars.There is some observational evidence suggestinga higher frequency for planets at smaller separa-tions around metal-rich stars which may supportsuch a metallicity–migration relationship (Jones2004; Sozzetti 2004). As a result, the planet–metallicity correlation may (at least in part) bedue to a metallicity–migration relationship andmay not simply signify a metallicity–formation re-lationship.

Additional observations of extrasolar planetsare necessary to distinguish between the compet-ing modes of planet formation and the conditionswhich affect planet migration scenarios. For ex-ample, the core accretion model requires severalMyr to form Jovian planets in a disk (Inaba etal. 2003), while gravitational instabilities couldpotentially form these planets around very youngstars. Jovian planets are suspected around a few∼ 1 Myr young stars (e.g., CoKu Tau/4 – Forrestet al. 2004; GQ Lup – Neuhauser et al. 2005);however, clear detection and firm mass determina-tions from techniques such as radial velocity vari-ations remain elusive on such young stars. Tran-siting extrasolar planets offer the opportunity todetermine both the mass and radius of the planet,and hence the planet’s density. Such planets al-low us to better constrain the variety of extraso-lar planet properties and offer another potentialway to distinguish between planet formation sce-narios. For example, the high core mass derivedfor the transiting Saturnian mass planet in orbitaround HD 149026 (Sato et al. 2005) has beeninterpreted as strong support for the core accre-tion model. There are now ∼ 20 known transitingextrasolar planets (http://exoplanets.eu/). Sev-eral of these planets have unexpectedly large radii(see Figure 3 of Bakos et al. 2007 and discussiontherein), perhaps suggesting additional heating ofthese planets other than that from irradiation bythe star they orbit. Additional effort is needed tounderstand the structure of extrasolar planets andtheir origins.

Here, we report the discovery of the third tran-siting planet from the XO Project. The XOProject aims to find planets transiting stars suffi-ciently bright to enable interesting follow up stud-

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ies (McCullough et al. 2005). 1 This planetpresents some very interesting properties: it has amass of Mpsini= 13.02±0.64 MJ and is in a shortperiod (3.2 d), eccentric orbit (e = 0.260 ± 0.017)around the apparently single F5V star GSC 03727-01064. XO-3b has a radius of 1.95±0.16 RJ, mak-ing it substantially larger in both radius and massthan most of the other reported transiting plan-ets. In §2 we discuss the observations leading tothe discovery of XO-3b. In §3 we present our anal-ysis of these observations to determine the stellarand planetary properties. In §4 we give a discus-sion of XO-3b in the context of other extrasolarplanets and current planet formation models, and§5 summarizes our conclusions.

2. Observations

2.1. XO Project Photometry

McCullough et al. (2005) describe the instru-mentation, operation, analysis, and preliminaryresults of the XO Project. In summary, the XOobservatory monitored tens of thousands of bright(V < 12) stars twice every ten minutes on clearnights for more than 2 months per season of vis-ibility for each particular star, over the periodSeptember 2003 to September 2005. From ouranalysis of more than 3000 observations per star,we identified XO-3 (Figure 1) as one of dozensof stars with light curves suggestive of a transit-ing planet. With the XO cameras on Haleakala,we observed four transits or segments of tran-sits of XO-3 in 2003 and three in 2004, on Ju-lian dates 2452934, 2452937, 2452998, 2453001,2453301, 2453320, and 2453352. Table 1 providesa sample of the photometry for XO-3 from theXO cameras. The full table is available in the on-line edition. From the survey photometry of XO-3(Figure 2), which has a nominal standard devia-tion of 0.8% or 8 mmag per observation, we de-termined a preliminary light curve and ephemeris,which we used to schedule observations of higherquality with other telescopes, as described in thenext subsections.

1This paper includes data taken on the Haleakala summitmaintained by the University of Hawaii, the Lowell Ob-servatory, the Hobby-Eberly Telescope, the McDonald Ob-servatory of The University of Texas at Austin, and fivebackyard observatories.

2.2. Additional Photometry

As outlined by McCullough and Burke (2006),once an interesting candidate is detedcted in theXO photometry, the candidate is released to anExtended Team (E.T.) of sophisticated amateurastronomers for additional observation. In 2006September through November and 2007 Januarythrough March, transit events of XO-3 were ob-served by members of the E.T. from a total of 4backyard observatories. Table 1 also provides E.T.photometry for XO-3. For the E.T. light curves,the median differential magnitude out of transitprovides the flux normalization and the standarddeviation out of transit provides the uncertaintyin the measurements. A fifth E.T. observatorywas used to obtain all-sky photometry for XO-3.These observatories are equipped with telescopesof aperture ∼ 0.3 m. The telescopes, equippedwith CCD detectors, are suitable for obtaininglight curves sufficient to confirm the nature of thetransit and obtain good timing information. Anetwork of such telescopes is well suited to observecandidates with known positions and ephemerides.Here, we use two such telescopes in North Amer-ica and two in Europe. Unlike the case of XO-1b(McCullough et al. 2006) and XO-2b (Burke et al.2007), we were not able to schedule observing timeon a larger (∼ 1-m) telescope to obtain additionalphotometry – the E.T. photometry is the best wehave obtained to date on XO-3b.

We estimate all-sky photometric B, V, RC, andIC magnitudes for XO-3 and several nearby refer-ence stars (Figure 1 and Table 2) calibrated usinga total of 6 Landolt areas (Landolt 1992). Usingthe Hereford Arizona Observatory (E.T. memberBLG) 0.36-meter telescope on photometric nights2006 October 27 and 2006 November 5, we mea-sured the fluxes of 48 and 18 Landolt stars, re-spectively at an air mass similar to that for XO-3and established the zero points of the instrumentalmagnitudes and transformation equations for thecolor corrections for each filter and the CCD. Thederived magnitudes for XO-3 differed on the twodates by 0.03 magnitudes or less in all 4 colors.The B, V, RC, and IC absolute photometric accu-racies are 0.03, 0.03, 0.03, and 0.05 mag r.m.s.,including both the formal error and an estimatedsystematic error. The Tycho magnitudes for XO-3listed in Table 3 transform (via Table 2 of Bessel2000) to Johnson V = 9.86, i.e. 0.06 mag (2-σ)

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fainter than our estimate.

2.3. Spectroscopy

In order to measure the orbital elements of thesystem, and in particular the mass ratio, as well asto determine the characteristics of the host star,we obtained spectra of XO-3 with two-dimensionalcross-dispersed echelle spectrographs (Tull et al.1995; Tull 1998) at the coude focus of the 2.7-mHarlan J. Smith (HJS) telescope and via a fiberoptic cable on the 11-m Hobby-Eberly Telescope(HET). Both telescopes are located at McDonaldObservatory. The HJS spectra were obtained ina traditionally scheduled manner, while the HETspectra were obtained in queue scheduled mode.An iodine gas cell is used on the HET to providethe wavelength reference for velocity determina-tion. While an iodine cell is available for HJS tele-scope, it is not a facility instrument and the obser-vations of XO-3 collected here were done as a partof a radial velocity survey of young stars (Huerta2007; Huerta et al. 2007) for which very high (fewm s−1) velocity precision is not required. Wave-length and resulting velocity calibration is accom-plished by taking thorium-argon reference lampspectra before and after each observation of XO-3at the HJS telescope. At the HJS we obtained onespectrum per night, and at the HET we obtainedor two spectra per night. At both telescopes,the spectral resolution was R ≡ λ/∆λ ≈ 60000,and data were obtained on a total of 21 nights.The two-dimensional echelle spectra were reducedusing IDL procedures described in Hinkle et al.(2000) which include bias subtraction, flat field-ing using a quartz lamp spectrum, and optimalextraction of the data. Table 4 gives a log of thespectral observations.

3. Analysis

3.1. Ephemeris

We follow the same procedure used by Burkeet al. (2007) for XO-2b to refine the ephemeris ofXO-3b. We adopt the transit midpoint calculatedfor the event of 2006 October 16 as our ephemeriszeropoint: 2454025.3967±0.0021. All measuredtransit midpoints used to refine the ephemeris arereported in Table 5. To determine the orbitalperiod, we minimize the χ2 difference betweenthe observed transit times and a constant-period

ephemeris model. Based upon 3 transit events ob-served either in multiple passbands by the sameobserver, or observed by different observers, we es-timate the 1−σ uncertainty in the time of the cen-ter of an individual transit to be 5 minutes, whichwe adopt as the uncertainty for each of the transitsobserved in a single bandpass by only one observer.The best-fit period is 3.1915426±0.00014days.

3.2. Radial Velocity Measurements

Using spectra obtained at the HET, we mea-sured XO-3’s radial velocities with respect to thetopocentric frame using iodine absorption lines su-perposed on the spectra of XO-3. We modeled theextracted spectra using a template stellar spec-trum and the absorption spectrum of the HET io-dine gas cell (Cochran 2000). For the analysis ofXO-1 (McCullough et al. 2006) and XO-2 (Burkeet al. 2007), a high resolution spectrum of theSun and the Earth’s atmosphere (Wallace et al.1998) was used for the template stellar spectrum.We followed this same procedure initially for XO-3; however, the resulting radial velocity uncertain-ties were substantially higher than we achieved forXO-1 and XO-2. We suspect that the higher vsiniof XO-3 compared to these other two stars con-tributes to this increased uncertainty, but we werealso concerned that the higher temperature of XO-3 relative to these stars (and the Sun) resulted inspectral differences large enough to increase theuncertainty further. Therefore, we repeated theHET radial velocity determinations using a highresolution (λ/δλ ∼ 60000) spectrum of the F5Vstar HD 30652 from the SPOCS sample of stars(Valenti & Fischer 2005, hereafter VF05) as thetemplate stellar spectrum. The effective temper-ature of HD 30652 is 6424 K (VF05), within 5K of our derived Teff for XO-3 (see below), andvsini = 16.8 km s−1 which is close to the value of18.54 we find for XO-3 below. The spectrum ofHD 30652 is very similar in appearance to that ofXO-3.

Using an IDL2 implementation of Nelder andMead’s (1965) downhill simplex χ2 minimizationalgorithm, “Amoeba,” we adjusted parameters ofour model spectrum to fit the observations. Themodel includes convolution of our model spectra

2IDL is a software product of ITT Visual Information Solu-tions.

4

with a best-fitting Voigt profile to approximatethe (slightly non-Gaussian) line-spread functionof the instrument. The free parameters of ourmodel are a continuum normalization factor, theradial velocity of the star, the radial velocity ofthe iodine lines (which represent instrumental de-viations from their expected zero velocity with re-spect to the observatory), and an exponent (op-tical depth scale factor) that scales the depths ofthe lines as an arbitrary method of adjusting thespectrum of HD 30652 to even more closely matchthat of XO-3. Due to the iodine absorption, wecould not estimate the continuum level by inter-polating between local maxima in the spectrum, soinstead we solved for the continuum iteratively, asrequired to improve the fit between our model andthe observations. In the manner described above,for each ∼15 A section of each individual spectrumwithin the region of the recorded spectrum withsignificant iodine absorption, 5100 − 5700 A, weestimated the radial velocity of the star. From theapproximately normal distribution of the resultingradial velocity estimates for each epoch, we calcu-lated the mean radial velocity and its uncertainty.The 1-σ internal errors of the radial velocity mea-surements from the HET spectra range from ∼ 120to ∼ 200 m s−1 per epoch. These uncertainties arean order of magnitude larger than those achievedon XO-1 (∼ 15 m s−1) and XO-2 (∼ 20 m s−1).We transformed our measured radial velocities tothe barycentric frame of the solar system and sub-tracted the mean radial velocity and report thesevalues in Table 4. These velocity measurementsare shown in Figure 4.

The HJS radial velocity measurement techniqueused here is described in detail in Huerta (2007)and Huerta et al. (2007). We summarize it here.The radial velocity shift of each observed spec-trum is determined by a cross correlation analysisof the observed spectrum with respect to a ref-erence spectrum. To avoid complications whichmight result from a spectral type mismatch, we useone particular HJS epoch (2454137.8215) of XO-3as the reference spectrum. Again, wavelength cal-ibration is done by averaging the wavelength so-lution from Thorium-Argon reference lamps takenbefore and after each stellar observation. A to-tal of 15 spectral orders which contain numerous,strong stellar lines and no detectable telluric ab-sorption lines are used in the cross correlation

analysis. The radial velocity shift from the 15orders are averaged to get the final radial veloc-ity shift for each observation, and the standarddeviation of this mean is computed and adoptedas the uncertainty of the cross correlation anal-ysis. Barycentric radial velocity corrections arethen applied to the observations to determine ac-curate relative velocities for all the HJS spectra ofXO-3.

In addition to the uncertainty found above fromthe order-to-order scatter in the HJS radial veloc-ity measurements, there is an uncertainty associ-ated with the fact that our wavelength calibrationis based on spectra (Thorium-Argon) that are notobserved simultaneously with the stellar spectrumas is the case for iodine cell observations. In or-der to evaluate this uncertainty, a total of 6 starsfrom the Lick Planet Search sample (Nidever etal. 2002; Butler et al. 1996; Cumming et al.1999) which are known to be stable at the 5 –20 m s−1 level are observed each night and ana-lyzed in exactly the same way as we treat the spec-tra of XO-3. The internal uncertainty in the ra-dial velocity shift measurements for these standardstars based on the measured order-to-order scat-ter is typically < 20 m s−1. For the 2007 Febru-ary observing run (observations 3-10 at the HJS),the time series of radial velocity measurements forthese standard stars show an average standard de-viation of 123 m s−1, with all but one of these starsshowing a measured standard deviation less than126 m s−1. Therefore, we adopt the average valueof 123 m s−1 as the uncertainty due to the non-simultaneous nature of the HJS wavelength cali-bration for the 2007 February HJS observationsof XO-3 and add it in quadrature to the uncer-tainty from the order-to-order scatter for each ob-servation. For the 2006 October observations, theradial velocity standard star data gives an uncer-tainty of 132 m s−1 due to the non-simultaneousnature of the target and wavelength calibrationspectra. In addition, Huerta (2007) and Huerta etal. (2007) show that there are systematic offsets ofthe order of 100 m s−1 in the radial velocity stan-dard star measurements from one observing run tothe next, and speculate that this is caused by re-placing the spectrometer slit plug between observ-ing runs (the slit plug is not moved during the ob-serving runs). Using the radial velocity standardstar data, we estimate that there is a systematic

5

shift of 140± 33 m s−1 between the 2006 Octoberand 2007 February observations. We correct our2006 October radial velocity measurements of XO-3 by this amount and add this associated uncer-tainty in quadrature with the other radial velocityuncertainties described above. The measured ra-dial velocity shifts and the total uncertainty arereported in Table 4 and shown in Figure 4 phasedto the ephemeris known from the transits. Theuncertainty for the radial velocity (= 0 m s−1) ofthe reference epoch (2454137.8215) is set equal tothe lowest uncertainty determined for the otherHJS relative radial velocity measurements.

An eccentricity approximately equal to zero isexpected theoretically for hot Jupiters in ∼ 3 dayorbits (Bodenheimer et al. 2001) and was our ex-pectation for XO-3b. Therefore, we expected si-nusoidal radial velocity variations for XO-3 with aphasing consistent with the transit observations.The first two observations obtained with the HJS,taken at phase ∼ 0.64 and ∼ 0.95 are inconsis-tent with this expectation, forcing us to considereccentric orbits. Many additional data points arerequired to adequately constrain an eccentric or-bit, so we performed intense observing of XO-3with both the HJS and the HET in late 2006 andearly 2007. When fitting the orbit, in additionto the radial velocity data, we also use the timeof mid transit as a constraint in the fitting. Ad-ditionally, since the HJS and HET data are ondifferent relative velocity scales, we treat as a freeparameter the offset between these two scales. Inthe orbit fitting, we keep as fixed the orbital periodand phases determined from the transit ephemeris,and treat as free parameters the center of massvelocity of the system, the eccentricity, the veloc-ity amplitude K for XO-3, the longitude of peri-astron, the phase of periastron passage, and theoffset between the HJS and HET radial velocitymeasurements. We use the nonlinear least squarestechnique of Marquardt (see Bevington & Robin-son 1992) to find the best fit parameters for theorbit, which are given in Table 6. Figure 4 showsthe radial velocity data along with the orbital so-lution. In this plot, the HET data are shifted (byadding 1192.52 m s−1 to the HET radial velocities)to match the HJS spectra and the center of massvelocity (1002.68 m s−1) of all points is subtractedfor display purposes, since our center of mass ve-locity is based largely on only relative data. (As a

result of these velocity shifts, the HJS velocities inTable 4 have had 1002.68 m s−1 subtracted fromthem for display in Figure 4, and the HET veloc-ities in Table 4 have had 189.84 m s−1 added tothem for display in Figure 4.) Uncertainties in theorbital fit parameters are derived by Monte Carlosimulation of the data: for 1000 simulations weconstruct fake radial velocity data using the or-bital fit and applying Gaussian random noise at alevel commensurate with the observed data, andthen fit this model data using the same procedureoutlined above. Because the orbital fit produces aminimum reduced χ2 = 0.53, we believe our radialvelocity uncertainty estimates may be a little toolarge. This can be seen in Table 4 where we reportthe observed minus calculated (O − C) velocitiesof the fit. These values are typically less than ourderived uncertainties for the radial velocity mea-surements. The standard deviation of the HJSO − C values is 122 m s−1 compared to a meanderived uncertainty of 156 m s−1. For the HETdata, the standard deviation of the O − C valuesis 90 m s−1compared to mean derived uncertaintyof 169 m s−1(though this is dominated by a sin-gle observation), indicating that the derived uncer-tainties for both data sets are too large. Reducingthese uncertainties would reduce the uncertaintiesin the fit parameters, but we leave them as given,since we can not justify any specific reduction inour measured radial velocity uncertainties. Fromthe fit, we find the velocity semi-amplitude of theorbit of XO-3 is K = 1471±48 m s−1.

As described below, the vsini of XO-3 is rela-tively large (18.5± 0.2 km s−1), which is substan-tially larger than most stars around which planetshave been discovered (HAT-P-2b is another ex-ample of a planet orbiting a star which has a rel-atively large vsini = 19.8 ± 1.6 km s−1, Bakos etal. 2007). The experience of Mandushev et al.(2005) and the similarity between the primary inthat system with XO-3 gives concern that XO-3could be a hierachical triple system in which twolower mass stars comprise a short period eclipsingbinary star which are themselves orbiting a moremassive, significantly brighter primary in a muchlonger period orbit. In the case of Mandushevet al. (2005), the short period binary stars dis-play tens of km s−1 velocity shifts, but the strong,broad lines of the system almost hide this signal.The apparent velocity shift of the system primary

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is then the result of weaker lines moving aroundthe (stationary) lines of the primary, producingline profile distortions which are misinterpreted asradial velocity variability of the primary.

Line bisector analysis is a standard techniqueto look for evidence that line distortions are pro-ducing false radial velocity variations (e.g., Hatzeset al. 1997; Martinez Fiorenzano et al. 2005), andhave been used successfully to identify false ra-dial velocity variations caused by star spots (e.g.,Queloz et al. 2001; Bouvier et al. 2007; Huertaet al. 2007) and hierarchical triple systems (e.g.,Santos et al. 2002; Mandushev et al. 2005). Toimprove signal-to-noise in the bisector analysis, itis common practice to compute the bisector of thespectrum cross correlation function (Queloz et al.2001; Santos et al. 2002; Mandushev et al. 2005).Here, we calculate the line bisector of the totalcross correlation function computed from all 15orders in the HJS spectra used in the determina-tion of the radial velocities. We then compute thebisector span, which is the difference of the line bi-sector at two reference levels. Here we define thebisector span, SB, as the value of the bisector at anabsolute cross correlation level of 0.15 minus thebisector at an absolute cross correlation level of0.75. These span values are given in Table 4. Be-cause the spectral regions of the HET data whichcontain strong stellar lines also contain numerousiodine absorption lines, we restrict ourselves to theHJS data for the analysis of line bisectors. Previ-ous studies which used bisector analysis to reveala false positive in planet searches find a clear cor-relation between the measured radial velocity andthe bisector span (Queloz et al. 2001; Santos et al.2002; Mandushev et al. 2005; Bouvier et al. 2007;Huerta et al. 2007). In these cases, whether dueto spots (e.g., Queloz et al. 2001; Bouvier et al.2007; Huerta et al. 2007) or due to a hirearchicaltriple system (Santos et al. 2002; Mandushev etal. 2005), the full range measured in SB is approx-imately the same as the full range of the measuredradial velocities.

On the other hand, Torres et al. (2004) use thelack of a correlation between SB and the radial ve-locity to help confirm a planet transiting OGLE-TR-56b. In addition, for OGLE-TR-56b the fullrange in SB is less than one third the full rangein the radial velocities, which Torres et al. (2004)argue is additional evidence for the planetary na-

ture of this object. Figure 5 shows the bisectorspan, SB, plotted versus the measured radial ve-locity for the HJS spectra of XO-3. Also given inthe figure is the value of the linear correlation co-efficient and its associated false alarm probability(Bevington & Robinson 1992) for these data. Inaddition to there being no significant correlationbetween SB and the radial velocity, the full rangeof SB is very small compared to the full range inthe radial velocity. Figure 5 and Table 4 show thatthe full range in SB in XO-3 is less than one tenththe full range in the radial velocity variations forXO-3.

As a last check that the photometric dimmingsand radial velocity variations of XO-3 are due toa planetary mass companion and not the result ofa hierarchical triple system with later type starspresent, we looked for an infrared (IR) excess inXO-3. O’Donovan et al. (2006) used a redderthan expected V − K measurement as one pieceof evidence to reject the planetary hypothesis forthe F star GSC 03885-00829. Using the photom-etry reported in Table 3, we estimate the colorV − Ks = 1.01 ± 0.06 for XO-3. Based on theeffective temperature and gravity found below forXO-3, we assign a spectral type of F5V basedon the calibrations presented in Cox (2000). Inthe color system of Bessell and Brett (1988), anF5V star such as XO-3 is expected to have anintrinsic (V −Ks)◦ = 1.10. For comparison to the2MASS colors reported in Table 3, we use the colortransformation relations found in the Explanatory

Supplement to the 2MASS All Sky Data Release

and Extended Mission Products by Cutri et al.(http://www.ipac.caltech.edu/2mass/releases/allsky/doc/),and find that we must add 0.039 mag to the Bessel& Brett value in order to estimate (V − Ks)◦ =1.14 on the 2MASS system. Thus, the observedV − Ks = 1.01 color of XO-3 is 0.13 mag bluer(∼ 2-σ) than its expected color. The slight bluingmay be the result of the relatively low metallicitywe derive (see below), but there certainly is noevidence for a near IR excess suggesting a hier-archical triple system. The bisector analysis andthe V − Ks color of XO-3 give us confidence thatboth the photometric dimming and radial velocityvariations we measure in this star are indeed dueto a planetary mass companion.

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3.3. Spectroscopically-Derived Stellar Prop-

erties and Planetary Mass

We used the software package SME (Valenti &Piskunov 1996) to fit each of the 10 spectra of XO-3 from the HJS telescope with synthetic spectra.We used the methodology of VF05, including theirminor corrections to match the Sun and removeabundance trends with temperature (negligible inthis case). Because of gaps between echelle or-ders, the HJS spectra are missing the wavelengthintervals 6000–6123 A, which was included in theValenti & Fischer analysis. These wavelength in-tervals are also missing from our extracted HETspectra because the relevant echelle orders spanthe gap between the two detectors.

We averaged our SME results for the 10 HJSspectra, obtaining the parameter values in Table7. Each value in the last column of the table,labeled “Precision” because systematic uncertain-ties are not included, is the standard deviationof the 10 measurements divided by

√9 to yield

the formal uncertainty in the mean. The medianvalue of each derived parameter (not given) dif-fers from the mean by less than the uncertainty inthe mean. The final row in the table gives [Si/Fe],which VF05 used as a proxy for alpha-element en-richment, when interpolating isochrones. Figure 6shows XO-3’s spectrum in the region of the Mg I Btriplet, which is the dominant spectroscopic con-straint on gravity. These three Mg lines also havea significant impact on the global [M/H] parame-ter, which is used to scale solar abundances for allelements other than Na, Si, Ti, Fe, and Ni.

Following the methodology of Fischer & Valenti(2005), we interpolated Yonsei-Yale (Y2) isochrones(Demarque et al. 2004) to determine probabil-ity distribution functions for the mass, radius,gravity, and age of XO-3. The trigonometricparallax of XO-3 is unknown, so we initially as-sumed distances of 240, 260, and 280 pc with anadopted uncertainty of 10 pc in each case (whichaffects the width of the resulting distribution func-tions). In order to perform the isochrone analysis,we need a V magnitude corrected for extinctionfor XO-3. We measured V = 9.80 ± 0.03 andB = 10.25 ± 0.03 for XO-3 (Table 3). These ob-servations give a measured B − V = 0.45 ± 0.04for XO-3. VandenBerg & Clem (2003) determinedempirical color-temperature relations for stars in-

cluding variations with gravity. We downloadedtheir high temperature table (mentioned in foot-note 1 of their paper) and interpolated in the tableto the effective temperature, surface gravity, andmetallicity of XO-3 determined above. Doing sogives a predicted (B − V )◦ = 0.440 for XO-3. Asdiscussed below, the lightcurve analysis and thederived properties of XO-3b favor a larger stellargravity for XO-3. So, we also interpolate the pre-dicted colors to a gravity of logg = 4.2 (a littlemore than 3σ higher than the spectroscopicallyderived value). Doing so results in a predicted(B − V )◦ = 0.446. These two predicted valuesare within 0.01 and 0.001 mag of the measuredB −V of XO-3, with an uncertainty of 0.04 in themeasured color. We therefore find no compellingevidence for significant reddening to XO-3 and weadopt V = 9.80± 0.03 as the extinction correctedmagnitude of the star.

We used our spectroscopic effective tempera-ture, spectroscopic gravity, and an assumed dis-tance to derive a bolometric correction by interpo-lating the “high temperature” table from Vanden-Berg & Clem (2003). We combined the bolometriccorrection with the observed V-band magnitudeto determine stellar luminosity. Then we used thestellar luminosity and our spectroscopic effectivetemperature, iron abundance, and alpha-elementenrichment to interpolate the Y2 isochrones toproduce the probability distribution functions inFigure 7.

The parameters we derive place XO-3 above thezero-age main sequence. As a result, reasonablefits to the data are possible for pre-main sequenceevolutionary tracks. However, if XO-3 were indeeda pre-main sequence star, we would expect strongLi I absorption at 6708 A. For example, in theirsurvey of Pleiades F stars, Boesgaard et al. (1988)show that the Li I line at 6708 A is approximatelyas strong as the nearby Ca I line at 6718 A isstars of similar spectral type to XO-3. We showin Figure 8 the Li line region in XO-3 produced byaveraging the spectra obtained at the HET (thisregion is free of iodine absorption). The positionof the Li I line is marked and the nearby Ca I lineis clearly visible. The Li line is not detectable,indicating that XO-3 is substantially older thanthe Pleiades, therefore we conclude XO-3 is not apre-main sequence star.

The most probable distance, mass, radius, and

8

age for XO-3 are 260 pc, 1.41 M⊙, 2.13 R⊙, and2.69 Gyr respectively, for our best fit gravity oflogg = 3.95. Combined with our measured semi-amplitude K = 1471±48 m s−1, this stellar massimplies a mass of Mpsini =13.02±0.64 MJ for theplanet XO-3b. If we assume that XO-3 is at 240 pcinstead, then the most probable mass, radius, age,and gravity are 1.36 M⊙, 1.95 R⊙, 2.78 Gyr, andlogg = 4.00, still placing the star well above themain sequence. As discussed below, one of the un-usual properties of XO-3b is the relatively large ra-dius we derive for the planet. The planetary radiusis approximately proportional to the inferred stel-lar radius. The planetary radius could be smallerif we could justify a smaller stellar radius; how-ever, doing so begins to bring the inferred gravityin conflict with the spectroscopically derived valueof logg = 3.95±0.062. Taking a 3σ upper limit, thelargest allowed value for the gravity is logg = 4.14,corresponding to a stellar mass of 1.27 M⊙ and astellar radius of 1.62 R⊙ for an inferred distance of200 pc. Together with the measured velocity semi-amplitude, these values imply a planetary mass ofMpsini = 12.14 ± 0.40 MJ where the uncertaintyhere is just that resulting from the radial velocityfit. Clearly, a trigonometric parallax measurementfor XO-3 will be of great value in better establish-ing both the stellar and planetary parameters.

3.4. Light Curve Modeling and the Plan-

etary Radius

For additional photometric analysis and lightcurvefitting, all the E.T. photometric data shown in Fig-ure 3 were averaged into 7.4 minute bins. Binningpermitted outlier rejection and empirical estima-tion of the noise by the scatter of the individualobservations, which is helpful in cases such as thisin which residual calibration errors can be signifi-cant. Differences in the depth and the shape of thetransit light curve in different photometric bandsare both expected and observed to be substantiallysmaller than the photometric uncertainties in theE.T. lightcurves, so we combined the light curvefrom all photometric bands to improve the signalto noise ratio. The binning begins with phasingthe light curves using the refined ephemeris de-termined in §3.1. The final binned light curveconsists of a robust average of E. T. photomet-ric measurements in bins 7.4 minutes in duration.This robust average is determined by first calcu-

lating the median and median absolute deviationfor each bin and rejecting points in the bin thatdeviate by more than 4σ from the median. Next,the mean and standard deviation are calculatedfor each bin and points which are more than 4σfrom the mean are rejected. This rejection basedon the mean and standard deviation is repeatedone more time before a final (“robust”) mean iscalculated. The final binned light curve is sup-plied in Table 1 and is shown in Figure 9 alongwith the fits described below. The uncertainty ineach binned light curve data point is the standarddeviation of (surviving) measurements in the bindivided by the square root of the number of mea-surements in that bin. The uncertainties in allbins are then multiplied by a correction factor asdescribed below (the uncertainties in Table 1 donot have this correction factor applied).

We modeled the mean transit light curve us-ing the analytic transit model of Mandel and Agol(2002). For a planet with an eccentric orbit, thetransit model has ten parameters: transit mid-point (to), orbital period, stellar mass (M⋆), stel-lar radius (R⋆), planet radius (Rp), inclination (i),quadratic limb darkening law coefficients (u1 andu2), eccentricity (e), and longitude of perihelion(ω). Throughout the following analysis, the or-bital period remains fixed at the period found in§3.1. The radial velocity data provide the bestestimates for e = 0.26 and ω = −15.4 that re-main fixed during the χ2 minimization. For eachfit presented below, we fix the stellar parametersto specific values based on the spectral synthesisand isochrone analysis presented above. We inter-polate quadratic limb darkening coefficients fromClaret (2000) for the Cousins R photometric band-pass (u1 = 0.220; u2 = 0.380) based on the grav-ity, effective temperature, and metallicity ([M/H])determined from the spectroscopic analysis (Ta-ble 7). The remaining parameters, to, Rp, andi, are allowed to vary and are solved for using the“Amoeba” algorithm mentioned earlier. For inputinto the Mandel and Agol (2002) transit model, wedetermine the projected separation between starand planet, δ, in the case of nonzero eccentric-ity following the procedure outlined by Hilditch(2001). For a given i, ω, and e, we determinethe true anomaly, θmin, that minimizes δ (Equa-tion 4.9 of Hilditch 2001). The resulting θmin cor-responds to the mean anomaly at transit midpoint

9

providing the zeropoint to convert observed timesto θ via Kepler’s equation.

We then initially fit the light curve assumingthe stellar mass (1.41 M⊙) and radius (2.13 R⊙)derived from the spectral fits as described above.The free parameters of the light curve fit are thenthe radius of the planet and the inclination of theorbit (and the time of mid-transit). The resultingfit is shown in the top panel of Figure 9. We thenrepeated the light curve fits using the extremesin the stellar mass (1.33 and 1.90 M⊙) and ra-dius (1.49 and 2.36 R⊙) which result from the 1σuncertainties in these parameters derived above.Together, this then gives us estimates of the plan-etary radius (Rp = 1.95 ± 0.16 RJ) and orbitalinclination (i = 79.◦32 ± 1.◦36) which are alsogiven in Table 6. The total mass of the planetis then 13.25±0.64 MJ. Examination of the figureshows that the model does not fit all aspects ofthe lightcurve: the model ingress and egress ap-pear noticeably longer than the observed ingressand egress.

In order to better model the shape of the ingressand egress while still maintaining the total dura-tion of the transit, the radius of the star and theplanet both need to be reduced and the inclinationmust increase so that the total chord length theplanet traverses accross the face of the star staysapproximately the same. Figure 7 shows that fromthe isochrone analysis, the stellar mass and radiusboth decrease as the distance is assumed to besmaller. We have performed the isochrone analysisevery 10 pc, so we computed model light curves,stepping down in distance 10 pc at a time fromour favored distance of 260 pc. For each new dis-tance, we adopt the corresponding stellar massand radius resulting from the isochrone analysisand compute χ2 for the fit to the light curve. Theminimum in χ2 occurs between 190 and 200 pc.Fitting a parabola to the lowest 4 χ2 values, givesa best fit to the light curve for a distance of 185pc which gives M∗ = 1.24 M⊙, R∗ = 1.48 R⊙, andlogg = 4.19. This fit is shown in the lower panelof Figure 9. As mentioned earlier, the the best χ2

is quite large (∼ 129) for the 63 degrees of free-dom in our fit, indicating that the photometricuncertainties have been underestimated. We at-tempt to correct this by determing the multiplica-tive factor required to give a χ2 = 63.0 (a reducedχ2 = 1.0). The uncertainties shown in Figure 9

have been multiplied by this factor. Having donethis, we can then assign as the uncertainty in RP

and i the difference between the best fit parame-ters and those we obtain when ∆χ2 = 1.0. Doingso, we find that the transit lightcurve is best fitfor Rp = 1.25 ± 0.15 RJ and i = 83.◦32 ± 1.◦26.The total mass of the planet is then 12.03 ± 0.46MJ. This value for the planetary radius is con-siderably lower than the value derived above fromthe default stellar parameters. As discussed be-low, we regard this as a lower limit to the trueplanetary radius; however, we note that the ambi-guity described here can be greatly diminished byobtaining very accurate photometry of a transit ofXO-3 and/or by obtaining a precise trigonometricparallax measurement.

4. Discussion

4.1. Comparison to other Transiting Plan-

ets

The ∼ 13 MJ planet XO-3b is unusual in manyrespects compared to the sample of known extra-solar planets. The mass of XO-3b is quite largecompared to most of the known extra-solar plan-ets. Zucker and Mazeh (2002) noted that the mostmassive short period planets are all found in mul-tiple star systems. Udry et al. (2003) emphasizedthe general lack of massive planets on short pe-riod orbits, particularly for planets orbiting sin-gle stars, and interpreted these results in terms ofplanetary migration scenarios. The star XO-3 isnot known to be a binary, though it is relativelyunstudied. While it is in the Tycho-2 catalog (Høget al. 2000), XO-3 does not have a significantproper motion measurement from the Hipparcos

mission. It is not known if any of the nearby starsseen in Figure 1 are common proper motion com-panions or not. Udry et al. (2003) point out thatamong planets orbiting single stars known at thattime, there were no planets more massive than 2MJ with periods less than 100 d (see also Eggen-berger et al. 2004). Using data for 218 extra-solar planets as compiled on the Geneva Extra-solar Planet website (http://www.exoplanets.eu)updated as of 27 May 2007, this general lack ofshort period, very massive planets is still appar-ent. There are only 3 other planets (HD 162020b,HD 17156b, and HAT-P-2b) orbiting apparentlysingle stars with periods less than 30 days and

10

masses larger than 2.5 MJ, and one of those (HD162020b) is suspected to actually be a much moremassive brown dwarf (Udry et al. 2002). There isonly 1 planet other than XO-3b with a mass largerthan 2.5 MJ and a period less than 4 days (HD120136b, M2sini = 4.14 MJ, Butler et al. 1997),and this star is known to be in a stellar binarysystem (Hale 1994).

The eccentricity of XO-3b is also quite rare rela-tive to other known extra-solar planets, given theshort orbital period of XO-3b. The eccentricityof the shortest period extra-solar planets are allquite low (e < 0.05) and most are consistent withzero eccentricity (Halbwachs et al. 2005). Thisis generally believed to be due to either tidal cir-cularization (e.g., Wu 2003; Ivanov & Papaloizou2007) or the result of smooth orbital migration ina disk which is not thought to produce significanteccentricity (Murray et al. 1998). The eccentric-ity of XO-3b is larger than all other planets withperiods less than that of XO-3b. We discuss theeccentricity of XO-3b further in Section 4.3. Whilethe mass, period, and eccentricity of XO-3b makeit quite rare among extra-solar planets, it is not to-tally unique. The recently discovered HAT-P-2b(Bakos et al. 2007) is very similar in many re-spects with a period of 5.6 d, Mp = 9.04 MJ, ande = 0.52. Also, the hot Neptune GJ 436b has aperiod of 2.6 days and e = 0.15 (Butler et al 2004;Deming et al 2007). Apparently the assumption,commonly-held prior to 2007, that short-periodplanets should have circular orbits was incorrect.

It is now well established that host star metal-licity correlates with the likelihood of finding aJupiter mass companion orbiting the star in thesense that higher metallicity stars are more likelyto have planets (Santos et al. 2004; Fischer &Valenti 2005). The metallicity of XO-3 ([M/H] =-0.20) is relatively low, and the above mentionedstudies show that planets are found around only∼ 3% of stars with similar metallicity. There areonly two other very massive planets (Mp > 5 MJ)known around stars (HD 111232, HD 114762) withmetallicities lower than that of XO-3. The correla-tion between host star metallicity and the proba-bility of finding a planet is often taken as supportfor the core accretion model (e.g., Pollack et al.1996; Bodenheimer et al. 2000; Hubickyj et al.2005) of planet formation. Under that assump-tion though, it may be surprising that a few very

massive planets have also been found around lowmetallicity stars. It has also been suggested thathost star metallicity decreases on average as theplanet mass increases which appears inconsistentwith the core accretion model (Ribas & Miralda–Escude 2007). On the other hand, gravitationalinstabilities in a disk are not expected to producea metallicity-planet correlation (Boss 2002). Tak-ing this notion further, Ribas and Miralda–Escude(2007) have suggested that the sample of extra-solar planets consists of objects formed via two dif-ferent paths: core accretion in a disk and fragmen-tation of a pre-stellar cloud. The high mass, rela-tively large eccentricity, and low metallicity of theXO-3b system would then “fit” the notion of Ribasand Miralda–Escude of a planet formed from thecollapse of a pre-stellar cloud, but of course theobservations of XO-3b do not prove (or disprove)either formation scenario.

4.2. How Big is XO-3b?

Potentially, one of the most unusual and inter-esting properties of XO-3b is the large radius ofthe planet. At an age of 2-3 Gyr, giant plan-ets and low mass brown dwarfs are expected tohave radii very close to that of Jupiter (Burrowset al. 2001; Baraffe et al. 2003). Using the ta-bles in Baraffe et al. (2003), the radius of a 13MJ planet is expected to be 1.03 RJ at 1 Gyr,and 0.97 RJ at 5 Gyr. Including heating fromthe central star can increase the expected radiusfor so-called “hot Jupiters” (e.g., Bodenheimer etal. 2003). Fortney et al. (2007) compute plane-tary structure models up to 11.3 MJ for a range oforbital separations, including the insolation pro-duced by absorption of solar radiation. At an ageof 300 Myr or older, all the models computed haveradii less than 1.3 RJ. In all cases analyzed byFortney et al., their models predict a substantiallysmaller radius than the 1.95±0.16 RJ determinedbased upon the values of the stellar mass and ra-dius inferred from the spectroscopic and isochroneanalysis of XO-3. Such a large discrepancy be-tween the observed and predicted planetary ra-dius is reminiscent of TReS-4 which has a radiusof 1.674±0.094 RJ (Mandushev et al. 2007). How-ever, the radius of XO-3b is quite uncertain, andthe light curve itself favors a smaller planetary ra-dius of Rp = 1.25± 0.15 RJ. Because XO-3b is inan eccentric orbit, for comparison with the models

11

of Fortney et al. (2007) we compute the averageorbital separation between XO-3b and its star andthen adjust this separation to take into accountthe difference in luminosity between XO-3 and theSun. The smaller planetary radius for XO-3b isassociated with a smaller stellar radius for XO-3,so for this exercise, we take R∗ = 1.57 R⊙. Wethus compute an effective distance of 0.024 AU touse for XO-3b when using the tables of Fortney etal. (2007). Interpolating in these tables, the pre-dicted radius for XO-3b is 1.18 RJ at 1 Gyr and1.11 RJ at 4.5 Gyr. Both values are lower thanthe smaller radius inferred for XO-3b, but only byabout 1σ. A more precise radius for XO-3b wouldbe a very interesting benchmark for the theory ofextra-solar planets.

The actual radius of XO-3b can be better con-strained by a true parallax measurement and moreprecise photometric observations (ideally in mul-tiple colors) of additional transits. The constraintfrom this work that favors the larger radius forXO-3b is the stellar gravity inferred from the spec-troscopic analysis. A number of investigators havepointed out that stellar gravity determinations arenotoriously difficult and the relative error in thisparameter is often substantially larger than almostall other measured properties of transiting extra-solar planetary systems. Many of these investiga-tors have elected to use the transit light curve fitsin combination with stellar isochrone models to es-timate the stellar and planetary radius and hencelogg for the star (e.g., Sozzetti et al. 2004, 2007;O’Donovan et al. 2006; Bakos et al. 2007; Man-dushev et al. 2007). The photometry of XO-3bis not as precise as that used in those studies, sowe do not emphasize the photometrically-derivedparameters as those investigators have. However,we can examine the spectroscopically determinedlogg to estimate by how much, and in what sense,it is in error.

The above studies typically use spectroscop-ically determined effective temperatures withisochrone models to determine the stellar grav-ity, implicitly assuming the gravities determinedfrom isochrone analyses are more accurate thanthe spectroscopic values. We can compare thegravities determined from the Y2 isochrones andour spectroscopic analysis for stars of similar spec-tral type to XO-3. The sample of stars studied byVF05 all have accurate Hipparcos parallax mea-

surements, allowing them to use the isochroneswith the spectroscopically determined Teff to de-termine the stellar gravity. This gravity can thenbe compared to the gravity VF05 derive from thespectra alone, using the same techniques used inthis paper. Figure 10 shows the result of thiscomparison for 79 stars with Teff > 6200 K fromVF05. Recall Teff = 6429 ± 50 for XO-3. Thereis considerable scatter in this figure; however, inthe immediate vicinity of the spectroscopic grav-ity determined for XO-3, the gravities determinedfrom the isochrones are lower. This is the op-posite sense to what is suggested by the transitlight curve of XO-3. As a result, there is no clearindication in this sample of stars that the spectro-scopically determined gravity for XO-3 is biasedlow. When making this comparison, it is appro-priate to consider the accuracy of the gravitiespredicted by the Y2 isochrones. Hillenbrand andWhite (2004) report excellent agreement (betterthat 3%) between dynamically determined stellarmasses from eclipsing binaries and masses deter-mined from the Y2 isochrones for main sequencestars more massive than ∼ 0.6 M⊙; therefore, wefind no reason to doubt the Y2 isochrones appro-priate for XO-3. We also note that results fromtransit light curve analyses are subject to variousassumptions and potential biases that may existin the models used for the fit. For example, Auf-denberg et al. (2005) confirm the prediction ofAllende Prieto et al. (2002) that 1D atmospheremodels, including the ATLAS models used byClaret (2000) to determine limb darkening coeffi-cients, predict too much limb darkening at opticalwavelengths, depending on the treatment of con-vection and convective overshoot in the model. Infitting light curves, such a bias will favor modelswith smaller impact parameters and smaller stellarradii. Therefore, until better data can be obtainedfor XO-3 (a precise parallax measurement and/ormore precise transit photometry), the stellar andhence planetary radius will remain significantlyuncertain. At this point, taking the photomet-ric lightcurve and the spectroscopic gravity intoaccount, we estimate the radius of XO-3b to be1.10 < Rp < 2.11 RJ including 1σ uncertaintieson both limits. The lower limit is almost certainlytoo low however, given the results of Aufdenberget al. (2005) and Allende Prieto et al. (2002) onlimb darkening showing that the values we have

12

used here are likely overestimated.

4.3. Tidal Circularization and the Eccen-

tricity of XO-3b

As mentioned in Section 4.1, an interesting as-pect of the orbit of XO-3b is its significant eccen-tricity, e = 0.260 ± 0.017, because most planetswith orbital periods similar to XO-3b are believedto have already been tidally circularized (Halb-wachs et al. 2005). Independent of exactly howXO-3b arrived at its current orbit, it is interest-ing to consider how long it can remain in such anorbit under the influence of tidal circularization.The question of tidal circularization of close extra-solar planets has been studied by several investi-gators (see Adams & Laughlin 2006 and referencestherein). The equations governing the tidal circu-larization of extra-solar planets depend on veryuncertain planetary quality factors, Qp (for ex-ample, see discussion in Gu et al. 2003) and insome cases equally uncertain stellar quality fac-tors, Q∗, if considering tides raised on the star bythe planet. As an illustration of the difficulty inestimating quality factors, Mathieu (1994) pointout that theoretically determined values of Q∗ im-ply a slow rate for close binary stars to circular-ize which is contradicted by observations in stellarclusters of various ages. As a result, many in-vestigators try to use the observations (of binarystars and extra-solar planets) to empirically esti-mate the quality factors. For Jupiter mass extra-solar planets, Qp ∼ 105−106 (e.g., Gu et al. 2003;Adams & Laughlin 2006).

Circularization timescales depend linearly onQp, so we assume Qp = 106 to get predictionsat the long end of what can currently be esti-mated. Using equation (18) of Gu et al. (2003)and equation (3) of Adams & Laughlin (2006)to estimate the circularization time for XO-3b,we find timescales of 0.29 Gyr and 0.33 Gyr, re-spectively, using values corresponding to the largeradius (1.9 RJ) for XO-3b. The circularizationtimescale in both studies depends on the recipro-cal of the planetary radius to the fifth power, sothe circularization timescale grows to 2.92 Gyr and3.32 Gyr, respectively, using values correspondingto the small radius (1.25 RJ) for XO-3b. Clearly,there is considerable uncertainty in the value of thecircularization time; however, because circumstel-lar disks appear to be lost after 10-20 Myr (e.g.,

Haisch et al. 2001; Mamajek et al. 2004), it ap-pears that XO-3b perhaps should have circularizedby now if it originally arrived at its current loca-tion while the circumstellar disk was still in place(i.e. through migration). If instead, XO-3b hasbeen scattered in to its current location (e.g., Ford& Rasio 2007) more recently, there may not havebeen enough time for circularization to occur, par-ticularly for the parameters corresponding to thesmaller radius for XO-3b. A third possibility isthat additional, undetected planets orbiting XO-3maintain the eccentricity of XO-3b (e.g., Adams& Laughlin 2006). It is interesting to note thoughthat for planets in short period, eccentric orbitssuch as XO-3b, the tides raised in the planet candeposit substantial energy into the planet (e.g.,Gu et al. 2003; Adams & Laughlin 2006) whichcan inflate it to radii coresponding to the largervalue we measure (Gu et al. 2003). It thereforeappears XO-3b can again serve as a very interest-ing benchmark for studies of tidal circularizationand tidal heating once a more accurate radius canbe established for the planet.

5. Summary

XO-3b is a massive planet or low mass browndwarf in a short period (3.2 d), eccentric (e = 0.26)orbit, around a somewhat evolved F5 star. Thereis relatively large uncertainty in the planetary pa-rameters owing to the unknown distance to XO-3.The mass of XO-3b is 11.57 – 13.97 MJ and theradius is 1.10 – 2.11 RJ. The larger mass and ra-dius are favored by our spectroscopic analysis ofXO-3, and the smaller values are favored by thelight curve analysis. A precise trigonometric par-allax measurement or more accurate photometriclight curve date are needed to distinguish betweenthese values.

We wish to thank the anonymous refereefor many useful comments which improved themanuscript. The University of Hawaii staffhave made the operation on Maui possible; wethank especially Bill Giebink, Les Hieda, JakeKamibayashi, Jeff Kuhn, Haosheng Lin, MikeMaberry, Daniel O’Gara, Joey Perreira, KailaRhoden, and the director of the IFA, Rolf-PeterKudritzki.

This research has made use of a Beowulf clus-

13

ter constructed by Frank Summers; the SIMBADdatabase, operated at CDS, Strasbourg, France;data products from the Two Micron All Sky Sur-vey (2MASS) and the Digitized Sky Survey (DSS);source code for transit light-curves (Mandel &Agol 2002); and community access to the Hobby-Eberly Telescope (HET), which is a joint projectof the University of Texas at Austin, the Penn-sylvania State University, Stanford University,Ludwig-Maximilians-Universitat Munchen, andGeorg-August-Universitat Gottingen. The HETis named in honor of its principal benefactors,William P. Hobby and Robert E. Eberly. Wethank the HET night-time and day-time supportstaff. We also wish to thank the support staff atMcDonald Observatory, and to especially thankM. Huerta for his assistance observing at the 2.7m HJS Telescope.

XO is funded primarily by NASA Origins grantNNG06GG92G and the Director’s DiscretionaryFund of STScI. C. M. Johns–Krull and L. Pratowish to ackowledge partial support from NASAOrigins of Solar Systems grant 05-SSO05-86. Wethank Brian Skiff and Josh Winn for noting that apreprint had an error in the coordinates and therewas a sign error in our radial velocity fitting algo-rithm, respectively; both were corrected prior topublication.

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16

Table 1

XO Survey & E.T. Light Curve Dataa

Heliocentric Julian Date Light Curve Uncertainty Filter Nb Observatory[mag] (1-σ) [mag]

2452932.10913 -0.0017 0.0027 Wc 1 XO2452932.10938 -0.0032 0.0026 Wc 1 XO2452932.11597 0.0024 0.0026 Wc 1 XO2452932.11621 0.0021 0.0025 Wc 1 XO2452932.12305 -0.0016 0.0026 Wc 1 XO

aThe complete version of this table is in the electronic edition of the Journal. Theprinted edition contains only a sample.

bAverage of N measurements

cThe filters used in the XO project telescopes are described in McCullough et al.(2005). The bandpass is essentially flat from 4000 – 7000 A.

Table 2

All-sky Photometric Magnitudes

Stara B V RC ICXO-3 10.25 9.80 9.54 9.28

1 11.37 10.82 10.50 10.182 12.38 11.62 11.17 10.763 14.81 13.07 12.17 11.304 10.13 9.03 8.43 7.885 12.67 11.87 11.43 10.926 9.26 8.78 8.50 8.217 15.52 13.46 12.41 11.508 15.34 13.84 12.59 11.51

aStars are identified in Figure 1.

17

Table 3

The Star XO-3

Parameter Value Reference

RA (J2000.0) 04h21m52s.71 a,bDec (J2000.0) +57◦49′01′′.9 a,bV, VT , Vcalc 9.80±0.03,9.904 ± 0.027,9.86 ± 0.027 c,b,dB − V, BT − VT , (B − V )calc 0.45±0.04,0.451 ± 0.040,0.383 ± 0.034 c,b,dV − RC 0.26±0.04 cRC − IC 0.26±0.06, cJ 9.013 ± 0.029 eJ − H 0.168 ± 0.034 eH − Ks 0.054 ± 0.026 eSpectral Type F5V cd 260 ± 23 pc c(µα, µδ) (−0.6 ± 2.7, 1.6 ± 2.6) mas yr−1 bGSC 03727-01064 a

References. —a) SIMBADb) Tycho-2 Catalogue, Høg et al (2000)c) this workd) Calculated from Tycho-2 measurements using standard transformationse) 2MASS, Skrutskie e al. (2006)

18

Table 4

Radial Velocity Shifts

Heliocentric Radial Velocity Uncertainty O − C BisectorJulian Date Shift [m s−1] (1 σ) [m s−1] [m s−1] Span [m s−1] Telescope

2454037.0112 1265 209 107 -31 HJS2454038.0098 2001 216 174 4 HJS2454137.8215 0 134 -92 0 HJS2454138.8124 216 134 -112 170 HJS2454139.8128 2625 139 -117 133 HJS2454140.8131 267 139 -47 203 HJS2454141.8047 77 149 13 112 HJS2454142.7967 2761 139 -1 68 HJS2454143.7999 881 143 242 42 HJS2454144.7967 -35 161 34 152 HJS

2454005.8587 1454 129 103 · · · HET2454006.8597 -976 122 16 · · · HET2454113.5842 -71 184 44 · · · HET2454121.7100 -981 188 -43 · · · HET2454121.7190 -965 204 -16 · · · HET2454122.5635 -1152 182 52 · · · HET2454127.7046 -240 195 32 · · · HET2454128.7048 -1271 183 21 · · · HET2454158.6220 1191 144 -32 · · · HET2454159.6091 -278 132 -31 · · · HET2454162.6193 -30 194 -248 · · · HET

19

Table 5

Transit Timing Measurements

Transit MidpointUT Date Observer Filter (HJD - 2450000.0)

2003 Dec 24 XO Wa 2997.727292004 Oct 22 XO Wa 3300.904792004 Nov 10 XO Wa 3320.057132004 Dec 12 XO Wa 3351.993412006 Sep 21 CF R 3999.860112006 Sep 21 MF R 3999.859132006 Oct 16 TV R 4025.396732006 Oct 20 EM V 4028.589842006 Oct 23 CF B 4031.777592006 Oct 23 CF R 4031.779542006 Oct 23 CF V 4031.780522006 Nov 08 CF B 4047.730222006 Nov 08 CF R 4047.738042006 Nov 08 CF V 4047.733402006 Nov 24 EM I 4063.690192006 Nov 24 MF I 4063.698972006 Nov 27 MF I 4066.885252007 Jan 26 EM I 4127.530762007 Feb 15 CF V 4146.677252007 Mar 19 MF R 4178.58936

aThe filters used in the XO project telescopes are de-scribed in McCullough et al. (2005). The bandpass isessentially flat from 4000 – 7000 A.

20

Table 6

Orbital Solution and Planetary Parameters XO-3b

Parameter Value Notes

P 3.1915426±0.00014 dtc 2454025.3967±0.0038 HJDe 0.260±0.017ω −15.4 ± 6.6 degreesT◦ 2454024.7278±0.0570 HJDK 1471±48 m s−1

Mpsini 13.02±0.64 bRp/Rs (0.92±0.04) × RJ/R⊙ ai 79.32±1.36 deg b,ca 0.0476±0.0005 A.U. bMp 13.25±0.64 MJ b,dRp 1.95±0.16 RJ a,b,c

Note.—

a) RJ = 71492 km, i.e. the equatorial radius ofJupiterb) for M∗ = 1.41±0.08 M⊙

c) for R∗ = 2.13±0.21 R⊙

d) MJ = 1.8988e27 kg

Table 7

Results of the SME Analysis

PrecisionParameter Mean (1 σ)

Teff [K] 6429 50logg [cm s−2] 3.95 0.062

V sin i [km s−1] 18.54 0.17[M/H] -0.204 0.023[Na/H] -0.346 0.046[Si/H] -0.171 0.017[Ti/H] -0.262 0.049[Fe/H] -0.177 0.027[Ni/H] -0.220 0.033[Si/Fe] 0.006 0.032

21

Table 8

Spectroscopically Derived Stellar parameters

Parameter @ 240 pc @ 260 pc @ 280 pc

1.33 1.36 1.39Mass [M⊙] 1.36 1.41 1.43

1.39 1.44 1.48

1.92 2.08 2.22Radius [R⊙] 1.95 2.13 2.27

1.99 2.17 2.34

3.98 3.93 3.87Log(g) [cm s−2] 4.00 3.95 3.89

4.02 3.97 3.91

2.66 2.53 2.41Age [Gyr] 2.78 2.69 2.68

3.06 2.83 2.96

Note.—For each parameter, the middle row is themaximum likelihood value, and the values in the rowsabove and below span the 68% likelihood of the prob-ability distributions (cf. Figure 7). The three columnscorrespond to three assumed distances for XO-3.

22

1

2

3

4

56

78

Fig. 1.— XO-3 is centered, indicated by the twohash marks. Stars from Table 2 are circled. Northis up; East to the left. The DSS image, digitizedfrom a POSSII-F plate with a IIIaF emulsion andan RG610 filter, subtends 15′ of declination.

0.0 0.5 1.0 1.5 2.0 2.5∆t (day)

0.98

0.99

1.00

1.01

Rel

ativ

e F

lux

−0.3 −0.2 −0.1 −0.0 0.1 0.2 0.3∆t (day)

0.98

0.99

1.00

1.01

Rel

ativ

e F

lux

Fig. 2.— A total of 2969 individual observationsof XO-3 by the two XO cameras over two seasons2004 and 2005 are shown wrapped and phased ac-cording to the transit ephemeris and averaged in0.01-day bins (line). The top panel is the fulllightcurve, and the bottom panel shows the re-gion around phase 0 (the primary transit, markedwith a vertical dashed line in the top panel) en-larged. From these data we identified the star as acandidate for more refined photometry with othertelescopes at epochs of expected transits (Figure3).

Fig. 3.— Time series photometry of XO-3 from2005 – 2007, with dates, observers, and filters in-dicated. The observations have been averaged in0.006-day bins. The figure is in color in the elec-tronic edition.

23

0.0 0.2 0.4 0.6 0.8 1.0Phase

−1500

−1000

−500

0

500

1000

1500

2000

Rad

ial V

eloc

ity S

hift

(m s

−1)

0.0 0.2 0.4 0.6 0.8 1.0Phase

−600−400

−200

0

200

400

600

RV

Res

idua

l (m

s−1

)

4000 4050 4100 4150HJD − 2450000.0 (days)

−600−400

−200

0

200

400

600

RV

Res

idua

l (m

s−1

)

Fig. 4.— Top: The measured radial velocity dataare shown along with the best fit orbit. The filledcircles are based on data from the 2.7 m Harlan J.Smith Telescope, and the filled squares come fromdata taken with the 11-m HET. The radial velocitycurve of XO-3 traces out an eccentric orbit with avelocity amplitude K = 1471±48 m s−1, implyingXO-3b’s mass is Mpsini= 13.02±0.64 MJ. The pe-riod and phase used to fold the measured radial ve-locities are fixed at values determined by the pho-tometric transits. The measured center of mass ve-locity with respect to the solar system’s barycenterhas been subtracted (its value is arbitrary since itis largely based on the intrinsically relative radialvelocity measurements from the HJS). Middle: Aplot of the radial velocity residuals, again phasefolded with the ephemeris determined from thephotometric transits. Bottom: The radial velocityresiduals are shown again, this time as a functionof time, showing no long term trend.

−500 0 500 1000 1500 2000 2500 3000Velocity (m s −1)

−50

0

50

100

150

200

250

Bis

ecto

r S

pan

(m s

−1)

r=−0.28, f=0.43

Fig. 5.— The bisector span is plotted as a func-tion of the measured radial velocity for all the HJSspectra of XO-3. There is no correlation of thespan with radial velocity and the full amplitude ofthe span variations is less than a tenth of the fullamplitude of the radial velocity variations.

5165 5170 5175 5180 5185 0.4

0.6

0.8

1.0

1.2 Fe Fe Fe Mg Fe Ni Fe Fe+ Fe Fe+ Fe Mg Ti Ni Fe Fe La+ Mg Ti+ Fe Ni Ti+ Fe Ti+ Ti

6018 6020 6022 6024 6026 6028

1.0

1.2 C Pr+ Fe Mn Pr+ Fe+ V Fe Fe Fe Fe Mn Fe Pr+ Ni Fe Fe V + Zr+ V + Ni Cr

6035 6040 6045

1.0

1.2 Mo V + Nd+ Fe Fe Nd+ Fe Fe Fe Ni V Si Mn S S Fe Ce+ Fe Fe+ Fe+ Zr S S S Pr+

6055 6060 6065

1.0

1.2 Fe Fe Ce+Nd+ S S Fe Cr+ Ni Fe Fe Fe Mn Sc+ Fe Fe Fe+ Cr Fe Fe Si Ti Fe Fe Si

6102 6104 6106 6108 6110 0.8

1.0

1.2 Fe Fe La+ Fe Fe Ca Fe Ca Fe Fe Fe+Sm+ Fe Ni Si Zr+ Zr+ Si Co Ni Nd+Ce+Sm+ Ni V

6124 6126 6128 6130 6132 6134 61360.8

1.0

1.2 Ni+ Si La+ Ti Zr Fe Fe Ni Co Cr+ La+ Fe Fe+ Ni Fe Si Si Nd+Nd+ Ni Zr V + V Fe Fe

6145 6150 6155

1.0

1.2 Zr Ce+ Si Fe Ti La+ Cr+ Fe+ Fe Pr+ Fe+ Ti Fe+ Fe Cr+ Na Si Si Ca O O Fe Nd+ Fe+ O

6165 6170 6175 Air Wavelength (Å)

0.8

1.0

1.2 Sm+ Na Fe+ Ca Ca Ni Fe Ca Fe Pr+ Ca Ca Fe Ca Fe Ni Ni La+ Fe Fe+ Ni Ni Ni Fe+Sm+

Res

idua

l Int

ensi

ty

Fig. 6.— The mean spectrum of XO-3 as observed(black histogram) and modeled with SME (curve,colored blue in the electronic edition) in the re-gion of the Mg b triplet. Labels note the elementsresponsible for the indicated spectral lines. Inter-mittent line segments (tan) beneath the horizon-tal axis indicate wavelength intervals used to con-strain the spectroscopic parameters. Very shortand intermittent line segments (black) immedi-ately above the spectrum indicate wavelength in-tervals used to constrain the continuum fit. Acolor version of this figure appears in the electronicversion of the paper.

25

150 200 250d (pc)

1.1

1.2

1.3

1.4

1.5

M/M

sun

1.4126

0 pc

150 200 250d (pc)

1.0

1.5

2.0

2.5

R/R

sun

2.12

150 200 250d (pc)

3.8

4.0

4.2

4.4

4.6

log g

3.95

150 200 250d (pc)

0.0

1.0

2.0

3.0

4.0

Age

(G

yr) 2.69

Fig. 7.— The stellar mass, radius, gravity, and agewhich result from the isochrone analysis. The solidline shows the derived values as a function of theassumed distance, and the gray region shows the68% confidence limit on these parameters. Dashedlines are drawn at 260 pc corresponding to thespectroscopically determined logg = 3.95.

6700 6705 6710 6715 6720Wavelength (Å)

0.85

0.90

0.95

1.00

1.05

Rel

ativ

e F

lux

Li I 6707.8 Å

Fig. 8.— The spectrum of XO-3 in the wavelengthregion of the Li I doublet at 6708 A. The posi-tion of the Li feature is marked. The Li featureis expected to be approximately as strong as theCa I line at 6718 A in Pleiades age stars of similarspectral type, and even stronger at younger ages.Therefore, we conclude XO-3 is not a pre-main se-quence star, but instead is likely on the post-mainsequence.

0.9920.9940.9960.9981.0001.0021.004

Rel

ativ

e F

lux

(a)

−0.15 −0.10 −0.05 −0.00 0.05 0.10 0.15∆t (day)

0.9920.9940.9960.9981.0001.0021.004

Rel

ativ

e F

lux

(b)

Fig. 9.— The combined E.T. transit light curveof XO-3 and the best fit transit light curve modelscomputed under different assumptions. The up-per panel (a) fixes the stellar mass and radius tothe values determined from the combined spectro-scopic and isochrone analysis. The lower panel (b)is determined by minimizing χ2 in the light curvefit by letting the stellar mass and radius vary ac-cording to the isochrone results for different as-sumed distances. This best fit produces a stellargravity (logg) more than 3σ greater than the mea-sured value from the spectroscopic analysis.

3.8 4.0 4.2 4.4 4.6 4.8 5.0 5.2log gSME

3.6

3.8

4.0

4.2

4.4

4.6

log

g Iso

XO

−3

[M/H] < 0.04 or > 0.04

79 VF05 stars with Teff > 6200

Fig. 10.— Comparison of spectroscopically de-rived gravities to those derived from isochroneanalysis for the sample of stars from VF05 whichhave effective temperatures similar to XO-3. Starswith [M/H]< 0.04 are shown as blue squares andthose with [M/H]> 0.04 are shown with red cir-cles. The green dashed line in Figure 10 is the lineof equality. The color version of the figure appearsin the electronic version of the paper.

26