K2-136: A Binary System in the Hyades Cluster Hosting a Neptune-sized Planet

EPIC 247589423 is a high proper motion star ($+81.8$ mas yr⁻¹ in right ascension and −35.2 mas yr⁻¹ in declination; UCAC4, Zacharias et al. 2013). In the 67 years since the 1950 Palomar Observatory Sky Survey (POSS) images, EPIC 247589423 has moved more than 6'', enabling us to utilize archival POSS data to search for background stars that are now, in 2017, hidden by EPIC 247589423. The Blue POSS1 image has better resolution ($\sim 2^{\prime\prime}$ versus $\sim 4^{\prime\prime}$) but the Red POSS1 image goes deeper (${\rm{\Delta }}\mathrm{mag}\sim 4$ versus ${\rm{\Delta }}\mathrm{mag}\sim 6$)

Using the 1950 POSS data (Figure 2), we find no evidence of a background star at the current position of EPIC 247589423 to a differential magnitude of ${\rm{\Delta }}B\sim 3$ mag in the blue and ${\rm{\Delta }}R\sim 4$ mag in the red. Because EPIC 247589423 is slightly saturated in the POSS images, this sensitivity was estimated by placing fake sources at the epoch 2017 position of EPIC 247589423 in the epoch 1950 images and estimating the 5σ threshold for detection. The photometric scale of the image (and hence, the magnitudes of the injected test stars) was set using the star located 40'' to the southeast of EPIC 247589423, which has an optical magnitude of approximately $B\approx 17$ mag and $R\approx 16$ mag.

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This analysis rules out a $\lesssim 10 \%$ eclipsing binary that is 3–4 mag fainter than the primary star, but it does not rule out the more extreme background eclipsing binaries (a 100% eclipsing binary could produce a 1500 ppm transit at a differential magnitude of ∼7 mag). However, this analysis was sufficient for us to initiate the remainder of the follow-up observations.

We performed both near-infrared and optical spectroscopy in order to characterize the host star properties and to search for secondary spectral lines.

3.2.1. IRTF SpeX

We observed EPIC 247589423 with the near-infrared cross-dispersed spectrograph SpeX (Rayner et al. 2003, 2004) on the 3 m NASA Infrared Telescope Facility on 2017 July 24 UT (Program 2017A019, PI C. Dressing). While available photometry indicates that the star is late-type, follow-up spectroscopy is essential to measure the spectral type and fundamental parameters.

We observed EPIC 247589423 under clear skies with an average seeing of ∼$0\buildrel{\prime\prime}\over{.} 7$. We used SpeX in its short cross-dispersed mode (SXD) with the 0.3 × 15${}^{{\prime\prime} }$ slit, allowing us to observe the star over $0.7\mbox{--}2.55\ \mu {\rm{m}}$ at resolution $R\sim 2000$. The target was observed at two locations along the slit in three AB nod pairs using a 50 s integration time in each frame, providing a total integration time of 300 s. The slit position angle was synced to the parallactic angle to avoid differential slit losses. An A0 standard, HD31411, was observed after our target and flat and arc lamp exposures were taken immediately after that, to allow for telluric correction and wavelength calibration using the data reduction package, SpeXTool (Vacca et al. 2003; Cushing et al. 2004).

SpeXTool performs flat fielding, bad pixel removal, wavelength calibration, sky subtraction, flux calibration, and spectral extraction and combination. The final extracted and combined spectra have signal-to-noise ratios (S/N) of 175 per resolution element in the J-band (1.25 μm), 217 per resolution element in the H-band (1.6 μm), and 208 per resolution element in the K-band (2.2 μm). The JHK-band spectra were compared to late-type standards from the IRTF Spectral Library (Rayner et al. 2009), seen in Figure 3. EPIC 247589423 is an approximate visual match to the K5 standard across all three bands. The increased noise visible in the regions of strong H₂O absorption is a result of increased telluric contamination, potentially due to the relatively large ∼19° separation between the primary target and the available A0 standard.

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Following the methods presented in Mann et al. (2013), we use our SpeX spectrum to estimate the fundamental parameters of effective temperature (${T}_{\mathrm{eff}}$), radius (${R}_{* }$), mass (${M}_{* }$), and luminosity (${L}_{* }$) for EPIC 247589423. We used the index-based temperature relations of Mann et al. (2013) to estimate the temperature in each of the J-, H-, and K-bands and calculated the mean of the three values. We estimated the uncertainty by adding in quadrature the standard deviation of the mean and the scatter in each of the Mann et al. (2013) index relations. The resulting ${T}_{\mathrm{eff}}$ = 4360 ± 206 K was then used to estimate the remaining stellar parameters and their uncertainties using the polynomial relations from Mann et al. (2013). We estimate ${R}_{* }/{R}_{\odot }$ = 0.674 ± 0.061, ${M}_{* }/{M}_{\odot }$ = 0.696 ± 0.070, and ${L}_{* }/{L}_{\odot }$ = 0.152 ± 0.052. The radius and mass yield an empirical stellar density of 3.2 ± 1.0 g cm⁻³.

We also used the TiO5 and CaH3 molecular indices from Lépine et al. (2003) to measure a spectral type of K7 ± 0.5. This visible index-based spectral type is consistent with the estimated stellar parameters (Pecaut & Mamajek 2013) and the visual comparison to standards in the near-IR. Given these results, we adopt a conservative dwarf spectral type of K5 ± 1 for EPIC 247589423. The late M type companion detected at close separation in near-IR adaptive optics imaging is ∼5–10 mag fainter than the K5 star at visible to near-IR wavelengths (see Section 3.3). Thus, it only contributes ≲1% of the flux across the wavelength ranges used in the SpeX analyses and does not significantly affect the results.

3.2.2. Keck HIRES

As part of our standard process for validating transiting exoplanets, we observed EPIC 247589423 with infrared high-resolution adaptive optics (AO) imaging, both at Keck Observatory and Palomar Observatory. The Keck Observatory observations were made with the NIRC2 instrument on Keck II behind the natural guide star AO system. The observations were made on 2017 August 20 in the narrow-band ${Br}-\gamma$ filter in the standard 3-point dither pattern that is used with NIRC2 to avoid the left lower quadrant of the detector, which is typically noisier than the other three quadrants. The dither pattern step size was $3^{\prime\prime}$ and was repeated three times, with each dither offset from the previous dither by $0\buildrel{\prime\prime}\over{.} 5$. The observations utilized an integration time of 3 s with one coadd per frame for a total of 27 s. The camera was in the narrow-angle mode with a full field of view of $10^{\prime\prime}$ and a pixel scale of approximately $0\buildrel{\prime\prime}\over{.} 1$ per pixel. The Keck AO observations clearly detected a faint companion approximately $0\buildrel{\prime\prime}\over{.} 7$ to the south of the primary target. However, good relative photometry of the detected companion was hampered by the fixed speckle pattern, which our post-processing was unable to fully remove.

EPIC 247589423 was re-observed with the $200^{\prime\prime}$ Hale Telescope at Palomar Observatory on 2017 September 06 utilizing the near-infrared AO system P3K and the infrared camera PHARO (Hayward et al. 2001). PHARO has a pixel scale of $0\buildrel{\prime\prime}\over{.} 025$ per pixel with a full field of view of approximately $25^{\prime\prime}$. The data were obtained with a narrow-band Br–γ filter $({\lambda }_{o}=2.166;{\rm{\Delta }}\lambda =0.02\,\mu {\rm{m}}),$ a narrow-band H-continuum filter $({\lambda }_{o}=1.668;{\rm{\Delta }}\lambda =0.0018\,\mu {\rm{m}}),$ and a standard J-band filter $({\lambda }_{o}=1.246;{\rm{\Delta }}\lambda =0.162\,\mu {\rm{m}}).$

The AO data were obtained in a 5-point quincunx dither pattern with each dither position separated by ${4}^{{\prime\prime} }$. Each dither position is observed three times with each pattern offset from the previous pattern by $0\buildrel{\prime\prime}\over{.} 5$ for a total of 15 frames. The integration time per frame was 4.2, 9.9, and 1.4 s in the Br–γ, H-cont, and J filters. We use the dithered images to remove sky background and dark current, and then align, flat-field, and stack the individual images. The PHARO AO data have a resolution of $0\buildrel{\prime\prime}\over{.} 10$ (FWHM) in the Br–γ filter and $0\buildrel{\prime\prime}\over{.} 08$ (FWHM) in the H-cont and J filters, respectively.

The sensitivities of the AO data were determined by injecting fake sources into the final combined images with separations from the primary targets in integer multiples of the central source's FWHM (Furlan et al. 2017). The sensitivity curves shown in Figure 4 represent the 5σ limits of the imaging data.

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The nearby stellar companion was detected in all three filters with PHARO. The companion separation was measured from the Br–γ image and found to be ${\rm{\Delta }}\alpha =0\buildrel{\prime\prime}\over{.} 10\pm 0\buildrel{\prime\prime}\over{.} 003$ and ${\rm{\Delta }}\delta =0\buildrel{\prime\prime}\over{.} 723\pm 0\buildrel{\prime\prime}\over{.} 03$. At the distance of the Hyades, the companion has a projected separation from the primary star of $\approx 40\,\mathrm{au}$. The AO imaging rules out the presence of any additional stars within ∼05 of the primary (∼30 au) and the presence of any brown dwarfs, or widely separated tertiary components beyond 05 ((∼30–1000 au). The presence of the blended companion must be taken into account to obtain the correct transit depth and planetary radius (Ciardi et al. 2015).

Table 1 presents the deblended magnitudes of both stars. The stars have blended 2MASS magnitudes of $J=9.343\pm 0.026$ mag, $H=8.496\pm 0.02$ mag and ${K}_{s}=9.196\pm 0.023$ mag. The stars have measured magnitude differences of ${\rm{\Delta }}J=4.97\,\pm 0.04$ mag, ${\rm{\Delta }}H=4.96\pm 0.03$ mag, and ${\rm{\Delta }}{K}_{s}=4.65\pm 0.03$ mag. Br–γ has a central wavelength that is sufficiently close to Ks to enable the deblending of the 2MASS magnitudes into the two components. The primary star has deblended real apparent magnitudes of ${J}_{1}=9.11\pm 0.04$ mag, ${H}_{1}=8.51\pm 0.02$ mag, and ${{Ks}}_{1}=8.38\pm 0.02$ mag, corresponding to ${(J-H)}_{1}\,=0.60\pm 0.05$ mag and ${(H-{K}_{s})}_{1}\,=\,0.13\pm 0.02$ mag; the companion star has deblended real apparent magnitudes of ${J}_{2}\,=14.1\pm 0.1$ mag, $H=13.47\pm 0.04$ mag, and ${{Ks}}_{2}=13.03\,\pm 0.03$ mag, corresponding to ${(J-H)}_{2}\,=\,0.63\pm 0.11$ mag and ${(H-{K}_{s})}_{2}\,=\,0.44\pm 0.05$ mag. Utilizing the $(\mathrm{Kepmag}\,-{Ks})\mathrm{versus}(J-{Ks})$ color relationships (Howell et al. 2012), we derive approximate deblended Kepler magnitudes of the two components of ${\mathrm{Kepmag}}_{1}=10.9\pm 0.1$ mag and ${\mathrm{Kepmag}}_{2}\,=17.4\,\pm 0.2$ mag, for Kepler magnitude difference of ${\rm{\Delta }}\mathrm{Kepmag}\,=6.5\pm 0.2$ mag, which is used when fitting the light curves and deriving a true transit depth.

The companion star has infrared colors that are consistent with M7/8V spectral type (Figure 5). It is unlikely that the star is a heavily reddened background star. Based upon an R = 3.1 extinction law, an early-F or late-A star would have to be attenuated by more than 6 mag of extinction to make the star appear as a late M-dwarf. The entire line-of-sight extinction through the Galaxy is only ${A}_{V}\approx 2$ mag (Schlafly & Finkbeiner 2011), making a background A or F star an unlikely source of the detected companion.

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The large 4'' pixels of Kepler mean we cannot isolate the transit to either the primary or secondary star, from the Kepler data alone. However, we rule out the possibility that the observed transits are of the M7/M8 star due to the lack of a secondary eclipse. We also show that the transit duration strongly favors a planet transiting the K5 star and not the M7/8 dwarf companion.

With a flux difference in the Kepler bandpass of 6.5 mag, the transit/eclipse would have to be $\gtrsim 65 \%$ deep in order to be occurring around the M7/8, after the dilution of the brighter K5V is taken into account. Such a deep eclipse would require the late-type star to be an eclipsing binary star system, which would require a secondary eclipse depth of $\lesssim 25 \%$. Even with the dilution of the primary star, a 25% eclipse would still produce an observed eclipse that is $\approx 1 \%$ deep. Yet, no secondary eclipse (to a limit of $\approx 0.03 \%$) is observed, indicating that the transit event is not a stellar eclipse around the companion M-star.

Additionally, the observed transit duration is more consistent with the orbiting event being around the primary K5V star rather than the M7/8V companion. The time between first and last contact is ${T}_{14}=3.59\pm 0.15\,\mathrm{hr}$. For a K5V star and the measured stellar radius of $R\sim 0.7$ ${R}_{\odot }$ and mass of $M\sim 0.7$ ${M}_{\odot }$, a circular orbit with a period of 17.3 days would have a transit duration of ${T}_{14}\approx 3.7\,\mathrm{hr}$. If instead the star that is transited is the M7/8V star, the stellar radius and mass reduce to $R\sim 0.1$ ${R}_{\odot }$ and $M\sim 0.1$ ${M}_{\odot }$, and corresponding transit duration would only last 1 hr—significantly shorter than the observed transit duration (Seager & Mallén-Ornelas 2003).

The transit duration could, of course, be longer if the orbit is not circular:

$$\begin{eqnarray}&&\displaystyle \frac{{t}_{\mathrm{ecc}}}{{t}_{\mathrm{circ}}}=\displaystyle \frac{\sqrt{(1-{e}^{2})}}{1+e\cos (\omega -90^\circ )}\end{eqnarray} \tag{ 1 }$$

where e is the eccentricity and ω is the argument of periastron (e.g., Kane et al. 2012). In order to achieve a transit duration near what is observed (∼3.5 hr), the eccentricity would need to be $e\gtrsim 0.85$ and the transit would need to occur near apoapsis. For any other argument of periastron, the eccentricity would need to be even larger. While this is not impossible, it seems a rather contrived scenario; of the 876 confirmed planets with eccentricity estimates, only 14 have a eccentricities of $e\gtrsim 0.8$ and only 7 confirmed planets have an eccentricity of $e\gtrsim 0.8$. Given that none of these systems have orbital periods less than 70 days, and these systems represent only 1%–2% of the 876 confirmed planets with measured eccentricities, we consider such a scenario for the planet presented here as unlikely.

Finally, the stellar density from the transit duration of the light curve is more consistent with the host star being a K5V star than being an M7/8V star. Based upon the Choi et al. (2016) models, a mid- to late-M-dwarf with a mass of $M\sim 0.1$ ${M}_{\odot }$ should have a stellar density near $\rho \gtrsim 30$ g cm⁻³. By comparison, the stellar density for a K5V star with a mass of $M\sim 0.7$ ${M}_{\odot }$ should be near $\rho \approx 3.5$ g cm⁻³, and this is in reasonable agreement with the derived stellar density, assuming a circular orbit (see Figure 6). We also note that the higher end of the measured stellar density distribution is most consistent with our adopted primary mass and radius.

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We also applied the vespa planet validation tool to this system. This tool assumes that a planet candidate orbits a single main-sequence star, so here we assume that the planet orbits the brighter of our two stars and that stars in the Hyades have converged onto the main sequence. In this analysis, which also incorporates our high-resolution imaging data and our exclusion of additional spectroscopic companions, vespa returns a false positive probability of 8 × 10⁻⁵. Because of the caveats already mentioned, we do not take this as the true false positive probability, but qualitatively it indicates that, if the planet orbits the brighter K5V star, then it is likely not a false positive.

We regard all these items—the lack of a secondary eclipse, the length of the transit duration, the agreement of the derived stellar density with that of a K5V star, and the vespa results—as sufficient evidence to indicate that the observed transit most likely occurs around the primary star, that it is caused by a planet, and that, given the transit depth and stellar radius (0.71${R}_{\odot }$), the transiting planet is Neptune-sized.

As in our team's previous work (Schlieder et al. 2016; Crossfield et al. 2017; Dressing et al. 2017b), we use the free BATMAN¹⁸ software (Kreidberg 2015) to derive transit parameters from our light curve. We ran light curve fits while imposing Gaussian priors on the limb-darkening coefficients, using values appropriate for stars of K5V stars and included the dilution of the transit caused by the blending of the K5V with the M7/8V star. The derived transit and planet parameters are presented in Table 2, and the final fit to the phase-folded light curve is shown in Figure 1(d). Based upon the transit fits and the HIRES stellar parameters, the transit is caused by a Neptune-sized planet (Rp $=\,{3.03}_{-0.47}^{+0.53}$ ${R}_{\oplus }$) orbiting the K5V primary star with an orbital period of $P\,=17.3077\pm 0.0013$ days.

Table 2.  Planet Parameters

Parameter	Symbol	Units	Value
Time of Transit Center	${T}_{0}-2454833$	BJDTDB	2997.0235 ± 0.0025
Orbital Period	P	days	17.3077 ± 0.0013
Orbital Inclination	i	deg	${89.30}_{-0.76}^{+0.49}$
Planet/Star Radius Ratio	${R}_{P}/{R}_{* }$	%	${3.85}_{-0.20}^{+0.47}$
Linear Limb Darkening	α	⋯	0.900 ± 0.030
Quadratic Limb Darkening	β	⋯	0.486 ± 0.030
Transit Duration (1st–4th)	T14	hr	${3.59}_{-0.14}^{+0.17}$
Transit Duration (2nd–3rd)	T23	hr	${3.22}_{-0.18}^{+0.15}$
Stellar Radius-orbit Ratio	${R}_{* }/a$	⋯	${0.0287}_{-0.0027}^{+0.0075}$
Impact Parameter	b	⋯	0.43 ± 0.28
Stellar Density	${\rho }_{* ,{circ}}$	g cm−3	${2.67}_{-1.34}^{+0.90}$
Semimajor Axis	a	au	0.11728 ± 0.00048
Planet Radius	RP	${R}_{\oplus }$	${3.03}_{-0.47}^{+0.53}$
Incident Flux	Sinc	${S}_{\oplus }$	${11.9}_{-3.2}^{+3.7}$
Secondary Eclipse Depth	${\delta }_{\mathrm{ecl}}\ (3\sigma )$	ppm	$\lt 238$

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We now refer to the planetary system as the following separate components: K2-136A is the primary K5V star and K2-136B is the M7/8V stellar companion. In the course of writing this paper, the authors became aware of a similar discovery paper (Mann et al. 2017b). That paper reports on the simultaneous discovery of the Neptune-sized planet reported here. The authors of that paper derive a very similar planetary radius (${R}_{{\rm{p}}}\approx 2.9\pm 0.1\ {R}_{\oplus }$ versus ${R}_{{\rm{p}}}\approx 3.0\pm 0.5\ {R}_{\oplus }$). In addition, they report two other planets in the system: an inner earth-sized planet (${R}_{{\rm{p}}}=0.99\pm 0.05\ {R}_{\oplus }$) and an outer super-Earth-sized planet (${R}_{{\rm{p}}}=1.45\pm 0.1\ {R}_{\oplus }$). In their paper, they "letter" the planets in order of orbital period: the inner planet as K2-136 b, the outer planet as K2-136 d, and the Neptune-sized planet, jointly discovered, is referred to as K2-136 c. We adopt the same lettering scheme in this paper; however, given our discovery of the stellar companion, the planets should should be referred to K2-136A b, K2-136A c, and K2-136A d.

The light curve is clearly modulated by stellar variability that appears to be quasi-periodic (Figure 1(b)). The full amplitude of the variations is $\sim 0.5 \%$ (∼5 mmag), which is comparable to field K-dwarfs (Ciardi et al. 2011). Being ∼6.5 mag fainter than the K-dwarf, the M-dwarf companion would need to have variability amplitudes on the order of 1–2 mag in order to produced the observed amplitude of variability. That level of variability associated with quasi-periodic rotation is typically not observed in the field or Hyades M dwarfs (Ciardi et al. 2011; Douglas et al. 2016). Thus, the (primary) source of the observed variability is likely the primary component of the system: K2-136A.

A Lomb–Scargle periodogram of the light curve shows its strongest peak at 15.2 ± 0.2 day, and an autocorrelation of the light curve shows its strongest (non-zero-lag) peak at 13.8 ± 1.0 day. Such a period is consistent with the periods of other Hyades members of a similar mass (Delorme et al. 2011; Douglas et al. 2016) and further evidence that the spot modulation pattern is due to the primary, rather than the secondary, which would be expected to be rotating more rapidly; it therefore seems possible that the stellar rotation period of K2-136A lies in this range.

A rotation period of 14–15 days is expected to produce an equatorial velocity for an 0.71 ${R}_{\odot }$ of ${V}_{\mathrm{eq}}\approx 2.4\mbox{--}2.5$ km s⁻¹. The HIRES spectrum yields $v\sin i$ = 3.9 ± 1.0 km s⁻¹, which is marginally consistent with the expected equatorial velocity derived from the rotation periods. The modest inconsistencies between our measured $v\sin i$ and expectations from the stellar radius and photometric rotation period might be accounted for by (1) systematic effects involved in our estimation of $v\sin i$, of order 1 km s⁻¹, and (2) surface differential rotation. Measured and expected differential rotation rates in K-dwarfs are in the range of ≲0.05 rad d⁻¹ (Barnes et al. 2005; Kitchatinov & Olemskoy 2012). If the modulation pattern in the K2 photometry is due to surface features at higher, more slowly rotating latitudes, it is possible the equatorial rotation period is shorter by ∼1 day.

While not definitive, the marginal agreement between the measured $v\sin i$ and the rotation period indicates that the star's rotational axis is nearly perpendicular to the orbital plane of K2-136A c. If there were a significant misalignment, we would expect a more significant difference between the light curve-derived rotation period and the measured $v\sin i$, although a longer time baseline would be useful to confirm this.

Finally, we note that the orbital period of the planet and the rotation period of the star are similar, but not the same. We estimated how long the planet would take to circularize (${\tau }_{\mathrm{circ}}$) using the equation given by Adams & Laughlin (2006):

$$\begin{eqnarray}&&{\tau }_{\mathrm{circ}}=1.6\ \mathrm{Gyr}\times \left(\displaystyle \frac{{Q}_{{\rm{p}}}}{{10}^{6}}\right)\times {\left(\displaystyle \frac{{M}_{* }}{{M}_{\odot }}\right)}^{-1.5}\\ &&\qquad \times \,\left(\displaystyle \frac{{M}_{{\rm{p}}}}{{M}_{\mathrm{Jup}}}\right)\times {\left(\displaystyle \frac{{R}_{{\rm{p}}}}{{R}_{\mathrm{Jup}}}\right)}^{-5}\times {\left(\displaystyle \frac{a}{0.05\mathrm{au}}\right)}^{6.5}.\end{eqnarray} \tag{ 2 }$$

The tidal circularization timescales linearly with Q_p, and the tidal parameter (Q_p) is notoriously uncertain. However, the Neptune value is estimated to be $\sim {10}^{5}$ with a possible range of ${10}^{4}\mbox{--}{10}^{6}$ (Maness et al. 2007), indicating that the circularization timescale for K2-136A c may be $\sim 500\mbox{--}600$ Myr—a timescale very similar to the age of the Hyades Cluster. If the tidal parameter is more akin to Jupiter (10⁶), the circularization timescale would be closer to 5 Gyr, well beyond the age of the Hyades.

K2-136A c is in a 17-day orbit around a K5V star and is experiencing a stellar insolation flux of S ∼ 12 ${S}_{\oplus }$ (Table 2), and as a result, it has an expected equilibrium temperature that is ${T}_{\mathrm{eq}}=400\mbox{--}600$ K depending on the planet albedo and the atmosphere re-circulation. The other previous detected transiting planet in the Hyades (K2-25b, Mann et al. 2016) orbits an M4.5V star, but has a much shorter orbital period of 3.485 days although it experiences a similar insolation flux (S ∼ 10 ${S}_{\oplus }$) as K2-136A c.

The two Neptune-sized planets in K2-25 and K2-136, respectively, occupy a similar location in the two-dimensional distribution of planet radius versus stellar insolation flux (see Figure 7). Both planets are at the edge of the distribution, perhaps indicating that the $\sim 0.6\mbox{--}0.8\,\mathrm{Gyr}$ Hyades planets, in comparison to the >Gyr Kepler sample used to define the evaporation valley, may still need to undergo significant evolution.

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If the known young cluster transiting planets are compared to the old field star transiting planets (primarily dominated by Kepler detections) in a two-dimensional distribution of planet radius versus stellar mass (see Figure 8), the distribution of the cluster planets does not look significantly different than the distribution of old field planets. This suggests that the long-term evolution of the cluster planets should lead them to the distribution of planets currently observed in the field.

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K2-136A c orbits a relatively bright star in the infrared ($K\sim 9$ mag) in a very well-studied and nearby open cluster and, thus, offers the opportunity for more detailed studies (e.g., with Spitzer and JWST). In the optical, the star is a bit fainter with $V\approx 11$ mag. If the K2-136A c has a similar density to Neptune, the expected radial velocity (RV) amplitude caused by the orbital motion of the planet should be on the order of 5 m s⁻¹, well within the reach of modern radial velocity spectrographs. As noted above, the system has been detected as an X-ray source, which may indicate that the star is active; however, this could be the M-dwarf companion and not the K-dwarf primary. The spectra of the K-dwarf appears to have an activity indicator of ${S}_{{HK}}=1.03$, suggesting that the RV jitter could be of 1–10 m s⁻¹ (Isaacson & Fischer 2010). Thus, an RV measurement of the planet's mass may be feasible. Such a measurement would provide an all-too-rare constraint on the bulk properties of a young sub-Neptune.

Unfortunately, the ecliptic latitude of K2-136A is $\sim 1^\circ$, and (like most K2 targets) it will not be observed in the prime TESS mission. However, the transit depth is approximately 1.5 mmag and, thus, ground-based observations of the transits may be possible to refine the transit ephemeris and to search for long-term timing variations indicative of the other planets in the system (e.g., Barros et al. 2017; Lendl et al. 2017).