EDP Sciences
Press Release
Free Access
Issue
A&A
Volume 561, January 2014
Article Number L9
Number of page(s) 7
Section Letters
DOI https://doi.org/10.1051/0004-6361/201322584
Published online 15 January 2014

© ESO, 2014

1. Introduction

In 2008 we began monitoring the radial velocities (RVs) of a sample of main sequence and giant stars in the open cluster (OC) M 67, to detect signatures of giant planets around their parent stars. An overview of the sample and of our first results is reported in Pasquini et al. (2012). The goal of this campaign is to study the formation of giant planets in OCs to understand whether a different environment, such as a rich cluster like M 67, might affect the planet formation process, the frequency, and the evolution of planetary systems with respect to field stars. In addition, searching for planets in OCs enables us to study the dependence of planet formation on stellar mass and to compare the chemical composition of stars with and without planets in detail. Stars in OCs share age and chemical composition (Randich et al. 2005), therefore it is possible to strictly control the sample and to limit the space of parameters in a better way than when studying field stars. To address these questions we started a search for planets around stars of the OC M 67. This cluster has solar age (3.5–4.8 Gyr; Yadav et al. 2008) and solar metallicity (+ 0.03 ± 0.01 dex; Randich et al. 2006). In this Letter, we present the RV data obtained for the stars YBP1194, YBP1514, and S364 that reveal the presence of Jovian-mass companions.

2. Stellar characteristics

The three stars belong to the M 67 sample presented in Pasquini et al. (2012) with a proper motion membership probability higher than 60% according to Yadav et al. (2008) and Sanders (1977), see Appendix A. The basic stellar parameters (V, B − V, Teff, log g and [Fe/H]) with their uncertainties were adopted from the literature. Considering a distance modulus of 9.63 ± 0.05 (Pasquini et al. 2008) and a reddening of E(B − V) = 0.041 ± 0.004 (Taylor 2007), stellar masses and radii were estimated using the 4 Gyr theoretical isochrones from Pietrinferni et al. (2004) and Girardi et al. (2000). The parameters derived from isochrone fitting are comparable, within the errors, with the values adopted from the literature. The main characteristics of the three host stars are listed in Table 1. We note that the errors on these values do not include all potential systematics (see Appendix A).

Table 1

Stellar parameters of the three M 67 stars hosting planets.

YBP1194 is a G5V star, described by Pasquini et al. (2008) as one of the five best solar analogs in their sample. A detailed spectroscopic analysis (Önehag et al. 2011) has confirmed the star as one of the best-known solar-twins.

YBP1514 also is a G5 main sequence star. We adopted the atmospheric parameters obtained by Smolinski et al. (2011), who used spectroscopic and photometric data from the original Sloan Digital Sky Survey (SDSS-I) and its first extension (SDSS-II/SEGUE). These values are consistent, within the errors, with what has been found in previous work on the same data by Lee et al. (2008) and in the study of Pasquini et al. (2008).

S364 (MMJ6470) is an evolved K3 giant star. The stellar parameters, summarized in Table 1, are taken from Wu et al. (2011). We derived its mass and radius by isochrone fitting (Pietrinferni et al. 2004).

3. Radial velocities and orbital solutions

The RV measurements were obtained using the HARPS spectrograph (Mayor et al. 2003) at the ESO 3.6 m telescope in high-efficiency mode (with R = 90 000 and a spectral range of 378–691 nm); the SOPHIE spectrograph (Bouchy & SOPHIE Team 2006) at the OHP 1.93 m telescope in high-efficiency mode (with R = 40 000 and a spectral range of 387–694 nm), and the HRS spectrograph (Tull 1998) at the Hobby Eberly Telescope (with R = 60 000 and a wavelength range of 407.6–787.5 nm). In addition, we gathered RV data points for giant stars observed between 2003 and 2005 (Lovis & Mayor 2007) with the CORALIE spectrograph at the 1.2 m Euler Swiss telescope. HARPS and SOPHIE are provided with a similar automatic pipeline to extract the spectra from the detector images and to cross-correlate them with a G2-type mask obtained from the Sun spectra. Radial velocities are derived by fitting each resulting cross-correlation function (CCF) with a Gaussian (Baranne et al. 1996; Pepe et al. 2002). For the HRS, the RVs were computed using a series of dedicated routines based on IRAF and cross-correlating the spectra with a G2 star template (Cappetta et al. 2012). All the observations for each star were corrected to the zero point of HARPS, as explained in Pasquini et al. (2012), and were analyzed together. Two additional corrections were applied to the SOPHIE data, to take into account the modification of the fiber link in June 2011 (Perruchot et al. 2011) and the low S/N ratio of the observations. For the first, we calculated the offset between RV values of our stellar standard (HD 32923) before and after the change of the optical setup. For the second, we corrected our spectra using Eq. (1) in Santerne et al. (2012). We studied the RV variations of our target stars by computing the Lomb-Scargle periodogram (Scargle 1982; Horne & Baliunas 1986) and by using a Levenberg-Marquardt analysis (Wright & Howard 2009, RVLIN) to fit Keplerian orbits to the RV data. The orbital solutions were independently checked using the Yorbit program (Segransan et al., in prep.). For each case we verified that the RVs did not correlate with the bisector span of the CCF (calculated following Queloz et al. 2001) or with the FWHM of the CCF. All the RV data for each star are available in Appendix A.

thumbnail Fig. 1

Top: Lomb-Scargle periodogram for YBP1194. The dashed lines correspond to 5% and 1% false-alarm probabilities, calculated according to Horne & Baliunas (1986) and white noise simulations. Bottom: phased RV measurements and Keplerian best fit, best-fit residuals, and bisector variation for YBP1194. Black dots: HARPS measurements, red dots: SOPHIE measurements, green dots: HRS measurements.

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YBP1194 We have acquired 23 RV measurements since 2008. Fifteen were obtained with HARPS with a typical S/N of 10 (per pixel at 550 nm), leading to a mean measurement uncertainty of 13 m s-1 including calibration errors. Eight additional RV measurements were obtained with SOPHIE and HRS with mean measurement uncertainties of 9.0 m s-1 and 26.0 m s-1. A clear 6.9 day periodic signal can be seen in the periodogram (see Fig. 1 top) with its one-year and two-year aliases on both sides (at 6.7 d and 7.03 d). A single-Keplerian model was adjusted to the data (Fig. 1 bottom). The resulting orbital parameters for the planet candidate are reported in Table 2. The residuals’ dispersion is σ(O–C) = 11.55 m s-1, comparable with the mean measurement accuracy (~15 m s-1), and the periodogram of the residuals does not show significant power excess, although structures are present.

Table 2

Orbital parameters of the planetary companions.

YBP1514 Twenty-five RV measurements have been obtained for YBP1514 since 2009: 19 with HARPS, the others with HRS and SOPHIE. The typical S/N is ~10 and the measurement uncertainty is ~15 m s-1 for HARPS, ~25 m s-1 for HRS, and ~10 m s-1 for SOPHIE.

thumbnail Fig. 2

Top: Lomb-Scargle periodogram for YBP1514. Bottom: phased RV measurements and Keplerian best fit, best-fit residuals, and bisector variation for YBP1514. Same symbols as in Fig. 1.

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A significant peak is present in the periodogram at 5.11 days (Fig. 2 top), together with its one-year alias at 5.04 days. We fitted a single-planet Keplerian orbit corresponding to the period P = 5.11 days (Fig. 2 bottom). The orbital parameters resulting from this fit are listed in Table 2. Assuming a mass of 0.96 M for the host star, we computed a minimum mass for the companion of 0.40 ± 0.11 MJup. The residuals to the fitted orbit have a dispersion of σ(O–C) = 14.6 m s-1, within the mean measurement uncertainty, and show no significant periodicity.

S364 We collected 20 RV measurements of S364 in about four years with HARPS, HRS, and SOPHIE. The average RV uncertainty is ~3.0 m s-1 for HARPS, ~7.0 m s-1 for SOPHIE and ~20 m s-1 for HRS. Seven additional RV measurements were obtained with CORALIE between 2003 and 2005, with a mean measurement uncertainty of ~12 m s-1.

thumbnail Fig. 3

Top: Lomb-Scargle periodogram for S364. Bottom: phased RV measurements and Keplerian best fit, best-fit residuals, and bisector variation for S364. Same symbols as in Fig. 1.

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The periodogram of the observed data is shown in Fig. 3 (top) and indicates an excess of power at ≈121.7 days. The other clearly visible peak at 182 days is the one-year alias of the planetary signal at P = 121.7 days. It disappears in the periodogram of residuals, which no longer shows any signal. We fitted a single-planet Keplerian orbit to this signal (Fig. 3 bottom) and found an orbital solution whose parameters are reported in Table 2. The residuals to the fitted orbit show a level of variation of σ = 16.0 m s-1, higher than the estimated accuracy, but the periodogram of the residuals does not reveal significant peaks.

4. Discussion and prospects

We have presented new results from our planet-search campaign in the OC M 67. Our measurements reveal that Y1194, Y1514, and S364 host planets.

To rule out activity-related rotational modulation as the cause of the RV variations in our object data, we investigated chromospheric activity in these stars by measuring the variations of the core of Hα with respect to the continuum. The low S/N ratio of our observations does not provide sufficient signal in the region of the more sensitive Ca II H and K lines. We followed a method similar to the one described in Pasquini & Pallavicini (1991). All the targets exhibit a very low level of activity: S364 shows a variability in Hα of 2%, YBP1514 and YBP1194 of 3% without significant periodicity. In addition, the M 67 stars have a very low level of chromospheric activity (Pace & Pasquini (2004): erg cm-2 s-1 for M 67 compared with erg cm-2 s-1 for the Hyades), which is not compatible with generating the high RV variations we observe. Therefore, rotationally modulated RV variations for the dwarfs in M 67 are certainly not a concern. The remote possibility that these stars are short-period binaries seen pole-on can also be excluded, because they are very active, and will show enhanced Hα cores and strong X-ray emission, which has not been observed for these stars (van den Berg et al. 2004). The fact that these stars are of solar age and that our research is focused on finding giant planets with an expected RV variability of tens of m s-1 makes the contamination by activity irrelevant.

It is remarkable that Y1194 is one of the best-known solar twins. This star together with Y1514, S364, and the other M 67 targets will be suitable for a detailed differential abundance analysis to compare the chemical composition of stars with and without giant planets.

All the orbital solutions show nonzero eccentricity, but this is also common among planets found around field stars. Quinn et al. (2013) explained that hot-Jupiters in OCs with nonzero eccentric orbits and circularization time-scales tcirc longer than the system age, might provide an observational signature of the hot-Jupiter migration process via planet-planet scattering. We evaluated tcirc for the eccentric orbits of YBP1194 and YBP1514. Assuming a tidal quality factor 6 × 104 < Qp < 2 × 106, we calculated 409 Myr < tcirc < 13.6 Gyr for YBP1194 and 220 Myr < tcirc < 6.9 Gyr for YBP1514 (see Quinn et al. 2013 for details). Given the solar age of M 67 and the wide range of possible tcirc, reflecting the choice of the Qp and the estimation of the planetary radius, no firm conclusion can be drawn for the origin of the eccentric short-period orbits of these stars. Moreover, further investigations and more RV data are needed to better constrain the eccentricities of these objects (see Pont et al. 2011).

The planet around the giant S364 belongs to the low-populated region of periods between ~10 and ~200 days and is one of the shortest periods found around evolved stars.

When we examine the current distribution of the Jupiter-mass planets for RV surveys around FGK stars we find an exoplanet host-rate higher than 10% for planets with a period of up to a few years and 1.20 ± 0.38% at solar metallicity, for very close-in hot-Jupiters with a period shorter than ten days (Cumming et al. 2008; Mayor et al. 2011; Wright et al. 2012). This rate around field stars has been in contrast to the lack of detected planets in both open and globular cluster for several years. Before 2012, the detections were limited to a long-period giant planet around one of the Hyades clump giants (Sato et al. 2007) and to a substellar-mass object in NGC2423 (Lovis & Mayor 2007). No evidence of short-period giant planets has been presented in the study of Paulson et al. (2004) around main-sequence stars of the Hyades, or in several transit campaigns (Bramich et al. 2005; Mochejska et al. 2005, 2006; Pepper et al. 2008; Hartman et al. 2009). These triggered the hypothesis that the frequency of planet-hosting stars in clusters is lower than in the field. To explain the dichotomy between field and cluster stars, it has been suggested that the cluster environment might have a significant impact on the disk-mass distribution. Eisner et al. (2008), studying disks around stars in the Orion Nebula Cluster (ONC), proposed that most of these stars do not posses sufficient mass in the disk to form Jupiter-mass planets or to support an eventual inward migration. Other scenarios may be attributed to post-formation dynamics, in particular to the influence of close stellar encounters (Spurzem et al. 2009; Bonnell et al. 2001) or to tidal evolution of the hot-Jupiters (Debes & Jackson 2010) in the dense cluster environment. van Saders & Gaudi (2011), in contrast, found no evidence in support of a fundamental difference in the short-period planet population between clusters and field stars, and attributed the nondetection of planets in transit surveys to the inadequate number of stars surveyed. This seems to be confirmed by the recent results. Indeed, we can list the discovery of two hot-Jupiters in the Praesepe OC in 2012 (Quinn et al. 2012) and of two sub-Neptune planets in the cluster NGC 6811 as part of The Kepler Cluster Study (Meibom et al. 2013), the new announcement of a hot-Jupiter in the Hyades (Quinn et al. 2013) and now the detection in M 67 of three Jupiter-mass planets presented in this work. Quinn et al. (2012) obtained a lower limit on the hot-Jupiter frequency in Praesepe of 3.8%, which is consistent with that of field stars considering the enriched metallicity of this cluster. Meibom et al. (2013) have found the same properties and frequency of low-mass planets in OCs as around field stars. In our case, for short-period giant planets we derived a frequency of 2% (errors computed according to Gehrels 1986); which is slightly higher than the value for field stars. Adding giant planets with long

periods, the rate becomes 3.4%, but this fraction is a lower limit that will increase with the follow-up of some other candidates (see Pasquini et al. 2012), which reveal suggestive signals for additional planetary companions. If these were confirmed, the frequency of giant planets would rise to 13%, in agreement with the rate of giant planets found by Mayor et al. (2011) for field stars.

Acknowledgments

L.P. acknowledges the Visiting Researcher program of the CNPq Brazilian Agency, at the Fed. Univ. of Rio Grande do Norte, Brazil. RPS thanks ESO DGDF, the Hobby Eberly Telescope (HET) project, the PNPS and PNP of INSU – CNRS for allocating the observations. M.T.R. received support from PFB06 CATA (CONICYT). We are grateful to Gaspare Lo Curto and Michele Cappetta for the support in data reduction analysis and helpful discussions.

References

Online material

Appendix A: CMD and membership probabilities

In this section we summarize the results presented in Pasquini et al. (2012), focusing in particular on the three stars discussed in the Letter.

YBP1194, YBP1514, and S364 belong to the M 67 sample that includes a total of 88 stars. All targets have V mag between 9 and 15, and a mass range between 0.9–1.4 M.

We selected main-sequence stars (included YBP1194 and YBP1514) with a membership probability higher than 60% and a proper motion shorter than 6 mas/yr with respect to the average according to Yadav et al. (2008). For the giants we refer to Sanders (1977). The RV membership was established for the latter following the work of Mermilliod & Mayor (2007), who studied the membership and binarity of 123 red giants in six old OCs, and of Mathieu et al. (1986), who made a very complete RV survey of the evolved stars of M 67 with a precision of a few hundreds of m s-1. The majority of the other stars were selected according to Pasquini et al. (2008), who used several VLT-FLAMES exposures for each star to classify suspected binaries. We found that YBP1194, YBP1514, and S364 are probable RV members with a mean RV within one-sigma from the average cluster RV. For the latter, we adopted the value of ⟨ RVM 67 ⟩ = 33.724 km s-1 and the dispersion of σ = ± 0.646 km s-1 estimated in Pasquini et al. (2012).

Table 3 shows proper motions and membership probability for the three stars discussed. Details about selection criteria and motion errors can be found in the original Yadav et al. (2008) and Sanders (1977) works.

In Fig. A.1, we report the observed region of the color-magnitude diagram (CMD), indicating in different colors the the position of the stars considered in this letter and the solar analog, as determined in Pasquini et al. (2008). The three stars analyzed in this work lie quite well on the cluster sequence in the CMD. We superimposed the isochrones from Pietrinferni et al. (2004) with solar metallicity and age corresponding to 3.5 Gyr (black curve), 4.0 Gyr (dark-blue curve) and 4.5 Gyr (light-blue curve). We also included the 4.0 Gyr isochrone (red curve) with a slightly lower reddening (E(B − V) = 0.02 instead of 0.041, Taylor 2007). This curve seems to match the colors of the turnoff better (see also the discussion in Pasquini et al. 2012). In the same figure, we report the Padova isochrone using E(B − V) = 0.041 ± 0.004, with solar metallicity, age 4.47 Gyr, and Y = 0.26 (Girardi et al. 2000).

Given that the values of stellar parameters have influence on the estimation of the planet masses, we evaluated the effects on the host star masses and radii of using isochrones with different ages and slightly lower reddening. While for the two

main-sequence stars YBP1194 and YBP1514 we found no significant incidence, for the giant S364, an age uncertainty of ±0.5 Gyr and a lower reddening would induce an error on the star mass of 4% and on its radii of 3%. Therefore, we decided to include this effect in the uncertainties of S364 listed in Table 1 and in the error of the planet mass.

thumbnail Fig. A.1

CMD of M 67 (photometry from Yadav et al. 2008) for probable members (Pμ > 60%). The isochrones are taken from the BaSTI library (Pietrinferni et al. 2004). The isochrones in black, dark blue, and light blue correspond to 3.5 Gyr, 4.0 Gyr, and 4.5 Gyr with a reddening E(B − V) = 0.041 ± 0.004 (Taylor 2007). The isochrone in red is a 4.0 Gyr with a lower reddening (E(B − V) = 0.02). The isochrone in orange is a 4.47 Gyr from Girardi et al. (2000) with E(B − V) = 0.041 ± 0.004. The location of the Sun, if it were within M 67, is marked with a ⊙ in yellow.

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Table A.1

Object ID, proper motions, and membership probability of the targets; reference.

Table A.2

Relative RV measurements, RV uncertainties, bisector span, and ratio of the Hα core with respect to the continuum (see Pasquini & Pallavicini 1991) for YBP1194.

Table A.3

Relative RV measurements, RV uncertainties, bisector span, and ratio of the Hα core with respect to the continuum (see Pasquini & Pallavicini 1991) for YBP1514.

Table A.4

Relative RV measurements, RV uncertainties, bisector span, and ratio of the Hα core with respect to the continuum (see Pasquini & Pallavicini 1991) for S364.

All Tables

Table 1

Stellar parameters of the three M 67 stars hosting planets.

Table 2

Orbital parameters of the planetary companions.

Table A.1

Object ID, proper motions, and membership probability of the targets; reference.

Table A.2

Relative RV measurements, RV uncertainties, bisector span, and ratio of the Hα core with respect to the continuum (see Pasquini & Pallavicini 1991) for YBP1194.

Table A.3

Relative RV measurements, RV uncertainties, bisector span, and ratio of the Hα core with respect to the continuum (see Pasquini & Pallavicini 1991) for YBP1514.

Table A.4

Relative RV measurements, RV uncertainties, bisector span, and ratio of the Hα core with respect to the continuum (see Pasquini & Pallavicini 1991) for S364.

All Figures

thumbnail Fig. 1

Top: Lomb-Scargle periodogram for YBP1194. The dashed lines correspond to 5% and 1% false-alarm probabilities, calculated according to Horne & Baliunas (1986) and white noise simulations. Bottom: phased RV measurements and Keplerian best fit, best-fit residuals, and bisector variation for YBP1194. Black dots: HARPS measurements, red dots: SOPHIE measurements, green dots: HRS measurements.

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In the text
thumbnail Fig. 2

Top: Lomb-Scargle periodogram for YBP1514. Bottom: phased RV measurements and Keplerian best fit, best-fit residuals, and bisector variation for YBP1514. Same symbols as in Fig. 1.

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In the text
thumbnail Fig. 3

Top: Lomb-Scargle periodogram for S364. Bottom: phased RV measurements and Keplerian best fit, best-fit residuals, and bisector variation for S364. Same symbols as in Fig. 1.

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In the text
thumbnail Fig. A.1

CMD of M 67 (photometry from Yadav et al. 2008) for probable members (Pμ > 60%). The isochrones are taken from the BaSTI library (Pietrinferni et al. 2004). The isochrones in black, dark blue, and light blue correspond to 3.5 Gyr, 4.0 Gyr, and 4.5 Gyr with a reddening E(B − V) = 0.041 ± 0.004 (Taylor 2007). The isochrone in red is a 4.0 Gyr with a lower reddening (E(B − V) = 0.02). The isochrone in orange is a 4.47 Gyr from Girardi et al. (2000) with E(B − V) = 0.041 ± 0.004. The location of the Sun, if it were within M 67, is marked with a ⊙ in yellow.

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In the text

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