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Issue
A&A
Volume 620, December 2018
Article Number A77
Number of page(s) 13
Section Planets and planetary systems
DOI https://doi.org/10.1051/0004-6361/201834073
Published online 29 November 2018

© ESO 2018

1 Introduction

Exoplanetary discovery has widened our perspective and knowledge of planetary science in the past two decades. The space-based mission Kepler used transit photometry to detect and characterise exoplanets (Borucki et al. 2010, 2011; Koch et al. 2010) and one of its key objectives is the determination of the frequency of terrestrial planets in the habitable zones of stars. From their sample of over 4000 transiting planet candidates, it was revealed that small planets (RP < 4.0 R⊕) are by far the most common in our Galaxy (Howard et al. 2012; Batalha et al. 2013; Dressing & Charbonneau 2013; Petigura et al. 2013); this result is also supported by radial velocity (RV) surveys (e.g. Mayor et al. 2011; Bonfils et al. 2013). While the Kepler sample provided an insight into the planet occurrence rate (e.g. Batalha 2014), only a few dozen host stars were bright enough for follow-up characterisation. With the loss of two reaction wheels on the Kepler spacecraft, the K2 mission was adopted to extend the transiting exoplanet discoveries (Howell et al. 2014). The K2 mission has observed 19 fields so far and has supplied precise photometry of approximately 20 000 bright stars per campaign. This has yielded hundreds of transiting planet candidates (e.g. Vanderburg et al. 2016; Barros et al. 2016; Pope et al. 2016), over 300 of which have been statistically validated (e.g. Montet et al. 2015; Barros et al. 2015; Crossfield et al. 2016).

Super-Earths are absent in our own solar system. Therefore, they are of particular interest in the study of planet formation and evolution. To probe the formation histories of these small planets, it is necessary to derive the planetary masses and radii with precision better than a few percent in order to differentiate their internal compositions in the context of planet evolution models (e.g. Zeng & Sasselov 2013; Brugger et al. 2017). Recent theories have proposed a distinct transition in the composition of small exoplanets (Weiss & Marcy 2014; Rogers 2015). Planets with RP ≲ 1.6 R⊕ typically have high densities and are predominantly rocky. On the other hand, planets with larger radii typically have lower densities and possess extended H/He envelopes. In fact, planets such as Kepler-10 b (RP = 1.42 ± 0.03 R⊕, ρP = 8.8 ± 2.5 g cm−3; Batalha et al. 2011), LHS1140 b (RP = 1.43 ± 0.10 R⊕, ρP = 12.5 ± 3.4 g cm−3; Dittmann et al. 2017), Kepler-20 b(RP = 1.87 ± 0.05 R⊕, ρP = 8.2 ± 1.4 g cm−3; Buchhave et al. 2016), and K2-38 b (RP = 1.55 ± 0.02 R⊕, ρP = 17.5 ± 7.35 g cm−3; Sinukoff et al. 2016) all have densities higher than that of the Earth (ρ⊕ = 5.5 g cm−3) and compositions consistent with a rocky world, whereas low density planets such as GJ 1214 b (RP = 2.68 ± 0.13 R⊕, ρP = 1.87 ± 0.40 g cm−3; Charbonneau et al. 2009), the Kepler-11 system (RP = 1.97–4.52 R⊕, ρP = 0.5–3.1 g cm−3; Lissauer et al. 2011), and HIP 116454 b (RP = 2.53 ± 0.18 R⊕, ρP = 4.17 ± 1.08 g cm−3; Vanderburg et al. 2015) have solid cores, and substantial gaseous envelopes.

Recent efforts by the California-Kepler Survey (CKS; Johnson et al. 2017; Fulton et al. 2017) have refined the physical characteristics of Kepler short-periodplanets (P < 100 days) and their host stars for an in-depth study of the planet size distribution. Their results show a significant lack of planets with sizes between 1.5 R⊕ and 2.0 R⊕. The gap in the radius distribution can be explained by the “photoevaporation” model (Owen & Wu 2013; Lopez & Fortney 2013), in which the gaseous envelopes of planets are stripped away as a result of exposure to high incident flux from their host stars. The CKS also highlighted the importance of obtaining precise measurements of planet masses and radii to perform statistically significant studies of the radius distribution.

In this paper, we report the detection of a 2.37-day transiting super-Earth, K2-265 b. Photometry by K2 and HARPS RV measurements were used to constrain the radius and mass measurements of this planet with a precision of 6 and 13%, respectively. In Sect. 2, we describe the observations made from K2, data reduction, and spectroscopic follow-up. Our analyses and results are presented in Sect. 3, and we conclude the paper with a summary and discussion in Sect. 4.

2 Observations

2.1 K2 photometry

K2-265 was observed during K2 Campaign 3 in long cadence mode. The photometry was obtained between November 2014 and January 2015. The target was independently flagged as a candidate from two transit searches; the first made use of the POLAR pipeline (Barros et al. 2016), and the second used the methods described in Armstrong et al. (2015a,b), where human input was involved to identify high priority candidates. K2-265 was also independently identified as a planet-hosting candidate by other search algorithms (Vanderburg et al. 2016; Crossfield et al. 2016; Mayo et al. 2018).

The K2 light curve generated from the POLAR pipeline (Barros et al. 2016) has less white noise than that of Armstrong et al. (2015a,b), hence the former was used in the planetary system analysis. The POLAR pipeline is summarised as follows: the K2 pixel data was downloaded from the Mikulski Archive for Space Telescopes (MAST)1. The photometric data was extracted using the adapted CoRoT imagette pipeline from Barros et al. (2014), which uses an optimal aperture for the photometric extraction. In this case, the optimal aperture was found to be close to circular and comprised of 44 pixels. The modified moment method developed by Stone (1989) was used to determine the centroid positions for systematic corrections. Flux and position variations of the star on the charge-coupled device (CCD) can lead to systematics in the data. These were corrected following the self-flat-fielding method similar to Vanderburg & Johnson (2014). Figs. 1 and 2 show the final extracted light curve and the phase-folded light curve of K2-265 respectively. Table 1 gives the photometric properties of K2-265.

Table 1

Properties of K2-265 have a nearby bound companion (see text for detailed description), hence values presented in this table are for the blended photometry.

2.2 Spectroscopic follow-up

We obtained RV measurements of K2-265 with the HARPS spectrograph (R ~110 000), mounted on the 3.6 m Telescope at ESO La Silla Observatory (Mayor et al. 2003). A total of 153 observations were made between 2016 October 29 and 2017 November 22 as part of the ESO-K2 large programme2. An exposure time of 1800 s was used for each observation, giving a signal-to-noise ratio of about 50 per pixel at 5500 Å. The data were reduced using the HARPS pipeline (Baranne et al. 1996). We computed the RV measurements with the weighted cross-correlation function (CCF) method using a G2V template (Baranne et al. 1996; Pepe et al. 2002), and we estimated the uncertainties in the RVs as described in Bouchy et al. (2001). The line bisector (BIS), and the full width half maximum (FWHM) were measured using the methods of Boisse et al. (2011) and Santerne et al. (2015). Ten observations that were obtained when the target was close to a bright Moon exhibit a significant anomaly in their FWHM, up to 500 ms−1. We removed these data completely from the analyses described in the later sections. The remaining 143 RV measurements and their associated uncertainties are reported in Table A.3. The time-series RVs and the phase-folded RVs of K2-265 are shown in Figs. 3 and 4, respectively. Following the calibrations of Noyes et al. (1984), we derived the activity index of . The activityindex is used in Sect. 3.3 to derive the stellar rotational period.

thumbnail Fig. 1

Detrended K2 transit light curve of K2-265 with positions of transits marked with blue dashed lines.
thumbnail Fig. 2

Phase-folded K2 light curve of K2-265 is shown in the top panel with the ephemeris from our analysis. The best-fit transit model is plotted in a red solid line and the residuals of the fit are plotted in the bottom panel.

2.3 Direct imaging observations

Shallow imaging observations were obtained with the NIRC2 instrument at Keck on 2015-08-04 in the narrow-band Brγ filter at 2.169 μm (programmeN151N2, PI: Ciardi). Several images were acquired with a dithering pattern on-sky and they were simply realigned and median-combined. In the combined image, a candidate companion was clearly detected at close separation from the star. Figure A.1 shows the K-bank Keck AO image of K2-265 and the nearby companion, where the contrast of the objects is measured to be Δ mag = 8.12 in the K band. The relative astrometry of the candidate was estimated using a simple Gaussian fitting on both the star and the candidate. The error on the measurement is conservatively estimated to ~0.5 pixel, i.e. ~5 mas. The relative Keck astrometry was derived following methods described in Vigan et al. (2016) and the following parameters were obtained: Δ α = −910 ± 5 mas, Δ δ = −363 ± 5 mas, separation = 979 ± 5 mas, and position angle = 248.27 ± 0.29 deg.

The target was further observed with the SPHERE instrument on the Very Large Telescope (VLT) instrument in the IRDIFS mode (Vigan et al. 2010; Zurlo et al. 2014). More details on these observations together with the data reduction are presented in Ligi et al. (2018). The relative astrometry of the candidate companion with respect to the star were derived from SPHERE/IRDIFS, and the results are shown in Table A.1. The combined astrometry confirms that the companion is bound with the target star. The SPHERE/IFS data was used to derive a low-resolution near-infrared (NIR) spectrum (Ligi et al. 2018), which we used to characterise the companion star and estimate its contamination in the K2 photometry (see Sect. 3.2).

2.4 Gaia astrometry

The Gaia Data Release 2 (DR2) has surveyed over one billion stars in the Galaxy (Gaia Collaboration 2016, 2018; Lindegren et al. 2018) and provided precise measurements of the parallaxes and proper motions for the sources. K2-265 has a measured parallax of 7.18 ± 0.05 mas, corresponding to a distance of 139 ± 1 pc. The proper motion of K2-265 is μRA = 30.20 ± 0.09 mas, μDec = −23.34 ± 0.06 mas. As part of the Gaia DR2, the stellar effective temperature of K2-265 was derived from the three photometric bands (Andrae et al. 2018) as K. The G-band extinction Ag = 0.101 and the reddening E(BP − RP) = 0.065 estimated from the parallax and magnitudes were used to determine the stellar luminosity, which in turn provides an estimate of the stellar radius as (Andrae et al. 2018). The stellar parameters from the results of Gaia DR2 are consistent with the distance estimate, effective temperature, and stellar radius, which are derived in the joint Bayesian analysis in Sect. 3.4. However, Gaia DR2 does not detect the companion star in the system and K2-265 is registered as a single object.

3 Analysis and results

3.1 Spectral analysis

The spectral analysis of the host star was performed by co-adding all the individual (Doppler corrected) spectra with IRAF3. We first derived the stellar parameters following the analysis of Sousa et al. (2008) by measuring the equivalent widths (EW) of Fe I and Fe II lines with version 2 of the ARES code4 (Sousa et al. 2015). We derived the chemical abundances using the 2014 version of the code MOOG (Sneden 1973), which used the iron excitation and ionisation balance. We obtained the following parameters: Teff = 5457 ± 29 K, log g = 4.42± 0.05 dex, [Fe/H] = 0.08 ± 0.02 dex, microturbulence ξt = 0.81 ± 0.05 km s−1. The errors provided in this work for the stellar parameters are precision errors that are intrinsic to the method (Sousa et al. 2011).

The chemical abundances of the host star are found in Table A.2. For more details on this analysis and the complete list of lines are provided in the following works: Adibekyan et al. (2012), Santos et al. (2015), and Delgado Mena et al. (2017). Li and S abundances were derived by spectral synthesis as performed in Delgado Mena et al. (2014) and Ecuvillon et al. (2004), respectively.

thumbnail Fig. 3

Time-series HARPS RV measurements (red circles) of K2-265. The best-fit Keplerian orbit of K2-265 b is plotted in black. The stellar activity is fitted with a GP. The grey region show the 1σ confidence interval of the GP.
thumbnail Fig. 4

Phase-folded HARPS RV measurements (black circles) of K2-265 as a function of the orbital phase. The black solid line is the best-fit RV curve. The binned RV measurements are denoted as red open circles.

3.2 Characterisation of the companion star

To determine the physical parameters of the bound companion, we used the same approach as in Santerne et al. (2016). We fit the magnitude difference between the target and companion star, as observed by SPHERE IRDIFS, with the BT-Settl stellar atmosphere models (Allard et al. 2012). The two stars are bound companions (see Sect. 2.3), hence they have the same distance to Earth and age, and they are assumed to have the same iron abundance. We used a Markov chain Monte Carlo (MCMC) method to derive the companion mass, using the results of the spectral analysis of the target star as priors on the analysis. We used the Dartmouth stellar evolution tracks to convert the companion mass (at a given age and metallicity) into spectroscopic parameters. Our final derivation gives Teff = 3428 ± 22 K, log g = 4.870 ± 0.017 (cgs), Mstar B = 0.40 ± 0.01 M⊙, , corresponding to a star of spectral type M2 (Cox 2000). Using this result, we integrated the SED models in the Kepler band and derived the contribution of flux contamination in the light curve of star A from star B to be 0.952 ± 0.024%. The derived contamination of the companion star was taken into account in the joint Bayesian analysis in Sect. 3.4 to determine the system parameters of K2-265. The parameters of the companion star and their corresponding uncertainties are reported in Table 2.

3.3 Stellar rotation

Rotational modulation is observed in the detrended K2 light curve as shown in Fig. 1. We derived the rotational period of K2-265 using multiple methods to determine the origin of the periodic variation.

We first calculated the stellar rotational period with the auto-correlation function (ACF) method as described in McQuillan et al. (2013, 2014) and found the stellar rotational period as 15.14 ± 0.38 d. We observed a further peak at 30.48 ± 0.28 d.

The Lomb–Scargle periodogram (Lomb 1976; Scargle 1982) analysis was performed to determine the periodicity in the RV data. Figure 5 shows the periodogram of the BIS analysis, FWHM, RV measurements, and S index. A clear peak is measured in all four periodograms at 32.2 ± 0.6 d, which is larger than but marginally consistent with the ACF period of 30.48 d at a 2σ level. The timescale of light-curve variation measures the changing visibility of starspots. We attribute the discrepancy between the two rotation periods to latitude variation of the magnetically active regions.

An upper limit of the sky-projected stellar rotational velocity was derived from the FWHM of the HARPS spectra (v sin i < 1.9 ± 0.2 km s−1). Using the stellar radius in Table 2, we estimate a rotation period Prot > 26.02 ± 3.08 d (assuming an aligned system, i = 90°), which agrees with the ~30 d period derived from the photometry and the RV data.

Furthermore, the stellar rotation period was also derived following the method of Mamajek & Hillenbrand (2008). In summary, we used the B-V colour from APASS to find the convective turnover time τc using calibrations from Noyes et al. (1984). We then used the measured Mount Wilson index SMW = 0.195 ± 0.025 to derive , from which we determine the Rossby number Ro = 1.94 using calibrations from Mamajek & Hillenbrand (2008). Finally, using the relation Prot = Ro × τc, we calculate the stellar rotation period as 32 ± 10 d.

Table 2

System parameters of K2-265 obtained from PASTIS.

3.4 JointBayesian analysis with PASTIS

We employed a Bayesian approach to derive the physical parameters of the host star and the planet. We jointly analysed the K2 photometriclight curve, the HARPS RV measurements and the spectral energy distribution (SED) observed by the APASS5, 2MASS, and the Wide-field Infrared Survey Explorer (WISE) surveys (Munari et al. 2014; Cutri 2014; a full list of host star magnitudes can be found in Table 1) using the PASTIS software (Díaz et al. 2014; Santerne et al. 2015). The light curve was modelled using the jktebop package (Southworth 2008) by taking an oversampling factor of 30 to account for the long integration time of the K2 data (Kipping 2010). The RVs were modelled with Keplerian orbits. Following similar approaches to Barros et al. (2017) and Santerne et al. (2018), we used a Gaussian process (GP) regression to model the activity signal of the star. We modelled the SED using the BT-Settl library of stellar atmosphere models (Allard et al. 2012).

We derived the system parameters via the MCMC method. The spectroscopic parameters of K2-265A were converted into physical stellar parameters using the Dartmouth evolution tracks (Dotter et al. 2008) at each step of the chain. The quadratic limb darkening coefficients were also computed using the stellar parameters and tables of Claret & Bloemen (2011).

For the stellar parameters, we used normal distribution priors centred on the values derived in our spectral analysis. We chose a normal prior for the orbital ephemeris centred on values found by the detection pipeline. Furthermore, we adopted a sine distribution for the inclination of the planet. Uninformative priors were used for the other parameters. The priors of the fitted parameters used in the model can be found in Table A.4.

Twenty MCMC chains of 3 × 105 iterations were run during the MCMC analysis, where the starting points were randomly drawn from the joint prior. The Kolmogorov–Smirnov test was used to test for convergence in each chain. We then removed the burn-in phase and merged the converged chains to derive the system parameters.

3.5 Stellar age

From the joint analysis of the observational data, together with the Dartmouth stellar evolution tracks, the age of K2-265 was determined as τiso = 9.7 ± 3.0 Gyr. The stellar rotation analysis in Sect. 3.3 found that K2-265 has a rotation period of ~ 30 d. We adopted a rotational period of 32.2 ± 0.6 d and followed the methods by Barnes (2010) to find that K2-265 has a gyrochronological age of τgyro = 5.34 ± 0.19 Gyr. We further derived the age of K2-265 using the relation between the [Y/Mg] abundance ratio and stellar age (Nissen 2015; Tucci Maia et al. 2016) and found an age of τ[Y/Mg] = 3.97 ± 2.59 Gyr. The value τ[Y/Mg] agrees with τgyro within 1σ uncertainty but is lower than the derived isochronal age. The low lithium abundance A(Li II) < 0.45 of the host star obtained from spectral analysis (Sect. 3.1) suggests that the host is not young. Hence, it is likely that the host is of at least an intermediate age.

4 Discussion and conclusions

K2-265 b has a mass of 6.54 ± 0.84 M⊕ and radius of 1.71 ± 0.11 R⊕ (Fig. 6). This corresponds to a bulk density of 7.1 ± 1.8 g cm−3, which is slightly higher than that of the density of the Earth. We applied a number of theoretical models to investigate the interior composition of the planet.

Fortney et al. (2007) modelled the radii of planets with a range of different masses at various compositions and derived an analytical function that allows an estimate of the rock mass fraction (rmf) of ice-rock-iron planets. We find a rmf of 0.84 for K2-265 b, which is equivalent to a rock-to-iron ratio of 0.84∕0.16, i.e. a rock fraction that is higher than the Earth. Seager et al. (2007) also used interior models of planets to study the mass–radius relation of solid planets. By assuming the planets are composed primarily of iron, silicates, water, and carbon compounds, Seager et al. (2007) showed that masses and radii of terrestrial planets follow a power law. Using the derived best-fit mass and radius of K2-265 b, the bulk composition of the planet was determined to be predominantly rocky with > 70% of silicate mantle by mass.

We performed a more detailed investigation of the composition of K2-265 b using the interior model of Brugger et al. (2017). This model considers planets composed of three differentiated layers: core (metals), mantle (rocks), and a liquid waterenvelope. Figure 7 shows the possible compositions of K2-265 b inferred from the 1σ uncertainties on the mass and radius of the planet. By focussing on terrestrial compositions only (i.e. without any water), we show that the central mass and radius of the planet are best fitted with a rock mass fraction of 81%, which is consistent with other theoretical predictions. However, given the uncertainties on the fundamental parameters, the rmf remains poorly constrained, namely, within the 44–100% range. If we assume that the stellar Fe/Si ratio (here 0.90± 0.41) can be usedas a proxy for the bulk planetary value (Dorn et al. 2015; Brugger et al. 2017), this range is reduced to 60–83%. In the case of a water-rich K2-265 b, the model only allows us to derive an upper limit on the water mass fraction (wmf) of the planet. Indeed, given the high equilibrium temperature of the planet (~ 1300 K assuming anEarth-like albedo), water would be in the gaseous and supercritical phases, which are less dense than the liquid phase. From the uncertainties on the mass, radius, and bulk Fe/Si ratio of K2-265 b, we infer that this planet cannot present a wmf larger than 31%.

The CKS measured precise stellar parameters of Kepler host stars using spectroscopic follow-up (Johnson et al. 2017)and refined the planetary radii to study the planet size distribution and planet occurrence rate (Fulton et al. 2017). The survey has revealed a bimodal distribution of small planet sizes. Planets tend to have radii of either ≈ 1.3 R⊕ or ≈ 2.4 R⊕, with a deficit of planets at ≈ 1.8 R⊕. The survey confirms the prediction by Owen & Wu (2013), whereby a gap in the planetary radius distribution exists as a consequence of atmospheric erosion by the photoevaporation mechanism. Alternatively, the core-powered mass loss mechanism could also drive the evaporation of small planets (Ginzburg et al. 2016, 2018).

Because of its close proximity to the host star, the super-Earth K2-265 b is exposed to strong stellar irradiation. The gas envelope of the planet could be evaporated as a result. This process was observed in a number of systems (e.g. HD209458 b; Vidal-Madjar et al. 2003, GJ 436 b; Ehrenreich et al. 2015). The present irradiance of the planet is S =S⊕(Ls∕L⊙)(AU/a)^2 = 9.32 × 105 W m−2, where Ls and L⊙ are the luminosity of the star and the Sun, S⊕ is the solar irradiance on Earth, and a is the semi-major axis of the planet. The equilibrium temperature of the K2-265 b can be estimated using Eq. (1) of López-Morales & Seager (2007): , where f and AB are the reradiation factor and the Bond albedo of the planet. Assuming an Earth-like Bond albedo AB = 0.3 and that the incident radiation is redistributed around the atmosphere (i.e. f = 1∕4), the equilibrium temperature of K2-265 b is Teff≈ 1300 K.

Indeed, K2-265 b lies below the lower limit of the photoevaporation valley as shown in the 2D radius distribution plot in Fig. 8. This implies that the planet could have been stripped bare from photoevaporation, revealing its naked core. This atmospheric stripping process is presumed to have occurred in the first ≈ 100 Myr since the birth of the planet when X-ray emission is saturated (Jackson et al. 2012), after which the X-ray emission decays. We estimated the total X-ray luminosity of K2-265 over its lifetime, , using the X-ray-age relation of Jackson et al. (2012). Using the results of Sect. 3.5, we adopted a mean age of 6.32 Gyr for the host star. The X-ray-to-bolometric luminosity ratio in the saturated regime for a B − V = 0.743 star is log (Lx∕Lbol) = −3.71 ± 0.05 ± 0.47. The corresponding turn-off age is logτsat = 8.03 ± 0.06 ± 0.31, where the decrease in X-ray emission follows a power law (α = 1.28 ± 0.17). Over the lifetime of the star, ergs (assuming efficiency factor η = 0.25) and K2-265 b is expected to have lost 2.7% of its mass under the constant-density assumption. K2-265 b has a predominantly rocky interior as shown in Fig. 7. This indicates that the planet was likely formed inside the ice line and could have either migrated to its current orbital separation well before ≈ 100 Myr or accreted its mass locally (Owen & Wu 2017).

K2-265 b is among the denser super-Earths below the photoevaporation gap. In addition to photoevaporation, giant impact between super-Earths could drive mass loss in the planetary atmosphere. Super-Earths are thought to have formed via accretion in gas discs, followed by migration and eccentricity damping due to their interactions with the gas disc (e.g. Lee & Chiang 2015), leading to densely packed planetary systems. As the gas disc disperses, secular perturbation between planets excites theireccentricity, triggering giant impacts between the bodies before the system becomes stable (Cossou et al. 2014). Two planets of comparable sizes could collide at a velocity beyond the surface escape velocity (Agnor & Asphaug 2004; Marcus et al. 2009). The impact could lead to a reduction in the planet envelope-to-core-mass ratio, hence an increase in the mean density and alteration of the bulk composition of the planet (Liu et al. 2015; Inamdar & Schlichting 2016).

Discoveries of super-Earths have shown a diversity of small planets in the mass–radius diagram. Precise RV and photometric measurements with an accuracy of a few percent are necessary to put strong constraints on the planetary mass and radius and provide a precise bulk composition. The core composition of the planet can be derived as a result. In particular, the mass fraction of a planetary core can inform us of the formation and evolution history of the planet. K2-265 b has a precisely determined mass (13%) and radius (6%) and the composition of the planet is consistent with a rocky planet. Its small radius and short orbital period suggest that K2-265 b could have been photoevaporated to a bare rocky core. Its high rock-to-mass fraction implies a planet formation within the ice line. Studying planets with an exposed core could provide valuable insight to planet formation via the core accretion mechanism. The increasing sample of small planets will help distinguish planet origins, identify types of mass loss mechanism, and probe the efficiency of atmospheric evaporation processes.

thumbnail Fig. 5

Lomb–Scargle periodogram of (panel a) BIS span; (panel b) FWHM; (panel c) RV; (panel d) S index (SMW). The peak position is indicated by the purple line and corresponds to a period of 32.2 ± 0.6 d. The orbital period of K2-265 b is indicated by the yellow line, and the planet signal is only significant in the RV at the 1% FAPlevel.
thumbnail Fig. 6

Mass–radius diagram of confirmed Earth-sized planets with masses up to 20 M⊕. Data were taken from NASA Exoplanet archive5. The mass–radius relations were taken from Zeng et al. (2016). From top to bottom: black solid lines denotes a pure water, pure rock, and pure iron composition. The grey dashed lines between the solid lines represent mass–radius relations for water-rock and rock-iron composites. The red solid line indicates the lower limit of a planet radius as a result of collisional stripping (Marcus et al. 2010). K2-265 b has a composition consistent with a rocky terrestrial planet.
thumbnail Fig. 7

Ternary diagram showing the possible composition of K2-265 b. The thick black line indicates the allowed composition of the planet inferred from the central values of the mass and radius of the planet, and the dashed lines denote the deviations from this line allowed by the 1σ uncertainties on the fundamental parameters. The red line and area show the compositions compatible with the planetary bulk Fe/Si ratio derived for K2-265 b from the stellar value. Compositions of the Earth and Mercury are shown for reference.
thumbnail Fig. 8

Planet radius distribution as a function of orbital period. The grey circles denote the planet sample obtained from the CKS sample (Fulton et al. 2017). The blue dot-dashed line and the green dashed line indicate the peak of the bimodal distribution of the planet radius distribution, where planets tend to favour radii of ~ 1.3 R⊕ and ~ 2.4 R⊕ due to the photoevaporation mechanism. The red dotted line indicates the lower limit of the photoevaporation valley derived from Owen & Wu (2017).

Acknowledgements

We thank the anonymous referee for the helpful comments that improved the manuscript. D.J.A.B. acknowledges support by the UK Space Agency. D.J.A. gratefully acknowledges support from the STFC via an Ernest Rutherford Fellowship (ST/R00384X/1). This work was funded by FEDER – Fundo Europeu de Desenvolvimento Regional funds through the COMPETE 2020 – Programa Operacional Competitividade e Internacionalização (POCI), and by Portuguese funds through FCT – Fundação para a Ciência e a Tecnologia in the framework of the projects POCI-01-0145-FEDER-028953 and POCI-01-0145-FEDER-032113. N.S. and O.D. also acknowledge the support from FCT and FEDER through COMPETE2020 to grants UID/FIS/04434/2013 and POCI-01-0145-FEDER-007672, PTDC/FIS-AST/1526/2014 and POCI-01-0145-FEDER-016886, and PTDC/FIS-AST/7073/2014 and POCI-01-0145-FEDER-016880. S.G.S acknowledges support from FCT through Investigador FCT contract nr. IF/00028/2014/CP1215/CT0002. S.C.C.B. also acknowledges support from FCT through Investigador FCT contracts IF/01312/2014/CP1215/CT0004. E.D.M. acknowledges the support by the Investigador FCT contract IF/00849/2015/CP1273/CT0003. R.L. has received funding from the European Union’s Horizon 2020 research and innovation programme under the Marie Skłodowska-Curie grant agreement n.664931. F.F. acknowledges support from PLATO ASI-INAF contract n.2015-019-R0. This research was made possible through the use of the AAVSO Photometric All-Sky Survey (APASS), funded by the Robert Martin Ayers Sciences Fund. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. This publication makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration. This research has made use of NASA’s Astrophysics Data System Bibliographic Services. This paper includes data collected by the K2 mission. Funding for the K2 mission is provided by the NASA Science Mission directorate. This research has made use of the Exoplanet Follow-up Observation Program website, which is operated by the California Institute of Technology, under contract with the National Aeronautics and Space Administration under the Exoplanet Exploration Program. This work has made use of data from the European Space Agency (ESA) mission Gaia (https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and Analysis Consortium (DPAC; https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has been provided by national institutions, in particular the institutions participating in the Gaia Multilateral Agreement.

Appendix A Supplementary tables and figures

Table A.1

NIRC2 and IRDIS astrometry of the candidate companion.

Table A.2

Chemical abundances of the host star, relative to the Sun.

thumbnail Fig. A.1

K-band Keck AO image shows a companion at a separation of = 979 ± 5 mas.
Table A.3

Radial velocity data.

Table A.4

Parameters used in the analysis.

References

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2

Based on observations made with ESO Telescopes at the La Silla Paranal Observatory under programme ID 198.C-0169.

3

IRAF is distributed by National Optical Astronomy Observatories, operated by the Association of Universities for Research in Astronomy, Inc., under contract with the National Science Foundation, USA.

4

The ARES code can be downloaded at http://www.astro.up.pt/~sousasag/ares/

All Tables

Table 1

Properties of K2-265 have a nearby bound companion (see text for detailed description), hence values presented in this table are for the blended photometry.

In the text
Table 2

System parameters of K2-265 obtained from PASTIS.

In the text
Table A.1

NIRC2 and IRDIS astrometry of the candidate companion.

In the text
Table A.2

Chemical abundances of the host star, relative to the Sun.

In the text
Table A.3

Radial velocity data.

In the text
Table A.4

Parameters used in the analysis.

In the text

All Figures

thumbnail Fig. 1

Detrended K2 transit light curve of K2-265 with positions of transits marked with blue dashed lines.
In the text
thumbnail Fig. 2

Phase-folded K2 light curve of K2-265 is shown in the top panel with the ephemeris from our analysis. The best-fit transit model is plotted in a red solid line and the residuals of the fit are plotted in the bottom panel.
In the text
thumbnail Fig. 3

Time-series HARPS RV measurements (red circles) of K2-265. The best-fit Keplerian orbit of K2-265 b is plotted in black. The stellar activity is fitted with a GP. The grey region show the 1σ confidence interval of the GP.
In the text
thumbnail Fig. 4

Phase-folded HARPS RV measurements (black circles) of K2-265 as a function of the orbital phase. The black solid line is the best-fit RV curve. The binned RV measurements are denoted as red open circles.
In the text
thumbnail Fig. 5

Lomb–Scargle periodogram of (panel a) BIS span; (panel b) FWHM; (panel c) RV; (panel d) S index (SMW). The peak position is indicated by the purple line and corresponds to a period of 32.2 ± 0.6 d. The orbital period of K2-265 b is indicated by the yellow line, and the planet signal is only significant in the RV at the 1% FAPlevel.
In the text
thumbnail Fig. 6

Mass–radius diagram of confirmed Earth-sized planets with masses up to 20 M⊕. Data were taken from NASA Exoplanet archive5. The mass–radius relations were taken from Zeng et al. (2016). From top to bottom: black solid lines denotes a pure water, pure rock, and pure iron composition. The grey dashed lines between the solid lines represent mass–radius relations for water-rock and rock-iron composites. The red solid line indicates the lower limit of a planet radius as a result of collisional stripping (Marcus et al. 2010). K2-265 b has a composition consistent with a rocky terrestrial planet.
In the text
thumbnail Fig. 7

Ternary diagram showing the possible composition of K2-265 b. The thick black line indicates the allowed composition of the planet inferred from the central values of the mass and radius of the planet, and the dashed lines denote the deviations from this line allowed by the 1σ uncertainties on the fundamental parameters. The red line and area show the compositions compatible with the planetary bulk Fe/Si ratio derived for K2-265 b from the stellar value. Compositions of the Earth and Mercury are shown for reference.
In the text
thumbnail Fig. 8

Planet radius distribution as a function of orbital period. The grey circles denote the planet sample obtained from the CKS sample (Fulton et al. 2017). The blue dot-dashed line and the green dashed line indicate the peak of the bimodal distribution of the planet radius distribution, where planets tend to favour radii of ~ 1.3 R⊕ and ~ 2.4 R⊕ due to the photoevaporation mechanism. The red dotted line indicates the lower limit of the photoevaporation valley derived from Owen & Wu (2017).
In the text
thumbnail Fig. A.1

K-band Keck AO image shows a companion at a separation of = 979 ± 5 mas.
In the text

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