During the first billion years of their life exoplanet atmospheres are modified by different atmospheric escape phenomena that can strongly affect the shape and morphology of the exoplanet itself. These processes can be studied with Ly H andor He I triplet observations. We present highresolution spectroscopy observations from CARMENES and GIARPS checking for He I and H signals in 20 exoplanetary atmospheres V1298 Tau c K2100 b HD 63433 b HD 63433 c HD 73583 b HD 73583 c K277 b TOI2076 b TOI2048 b HD 235088 b TOI1807 b TOI1136 d TOI1268 b TOI1683 b TOI2018 b MASCARA2b WASP189 b TOI2046 b TOI1431 b and HATP57 b. We report two new highresolution spectroscopy He I detections for TOI1268 b and TOI2018 b and a H detection for TOI1136 d. Furthermore we detect hints of He I for HD 63433 b and H for HD 73583 b and c which need to be confirmed. The aim of the Measuring Outflows in Planets orbiting Young Stars MOPYS project is to understand the evaporating phenomena and test their predictions from the current observations. We compiled a list of 70 exoplanets with He I andor H observations from this work and the literature and we considered the He I and H results as proxy for atmospheric escape. Our principal results are that 0.11 Gyr planets do not exhibit more He I or H detections than older planets and evaporation signals are more frequent for planets orbiting 13 Gyr stars. We provide new constraints to the cosmic shoreline the empirical division between rocky planets and planets with atmosphere by using the evaporation detections and we explore the capabilities of a new dimensionless parameter RSUBHeSUBRSUBHillSUB to explain the He I triplet detections. Furthermore we present a statistically significant upper boundary for the He I triplet detections in the TSUBeqSUB versus SUBpSUB parameter space. Planets located above that boundary are unlikely to show He I absorption signals.
['techniques: photometric', 'techniques: radial velocities', 'planets and satellites: atmospheres', 'planets and satellites: gaseous planets', 'planets and satellites: physical evolution', 'Astrophysics - Earth and Planetary Astrophysics']
The MOPYS project A survey of 70 planets in search of extended He I and H atmospheres No evidence of enhanced evaporation in young planets
{'No evidence of enhanced evaporation in young planets': 'J. Orell-Miquel 1 , 2 , F. Murgas 1 , 2 , E. Pallé 1 , 2 , M. Mallorquín 1 , 2 , M. López-Puertas 3 , M. Lampón 3 , J. Sanz-Forcada 4 , L. Nortmann 5 , S. Czesla 6 , E. Nagel 5 , I. Ribas 7 , 8 , M. Stangret 9 , J. Livingston 10 , 11 , 12 , E. Knudstrup 13 , 14 , S. H. Albrecht 14 , I. Carleo 1 , 2 , J. A. Caballero 4 , F. Dai 15 , 16 , 17 , E. Esparza-Borges 1 , 2 , A. Fukui 18 , 1 , K. Heng 19 , Th. Henning 20 , T. Kagetani 21 , F. Lesjak 5 , J. P. de Leon 21 , D. Montes 22 , G. Morello 13 , 1 , N. Narita 18 , 10 , 1 , A. Quirrenbach 23 , P. J. Amado 3 , A. Reiners 5 , A. Schweitzer 24 , and J. I. Vico Linares 25  \n(A ffi liations can be found after the references)  \nReceived 31 January 2024 / Accepted 30 May 2024', 'ABSTRACT': 'During the first billion years of their life, exoplanet atmospheres are modified by di ff erent atmospheric escape phenomena that can strongly a ff ect the shape and morphology of the exoplanet itself. These processes can be studied with Ly α , H α , and / or He i triplet observations. We present high-resolution spectroscopy observations from CARMENES and GIARPS checking for He i and H α signals in 20 exoplanetary atmospheres: V1298Tauc, K2-100b, HD63433b, HD63433c, HD73583b, HD73583c, K2-77b, TOI-2076b, TOI-2048 b, HD 235088 b, TOI-1807 b, TOI1136 d, TOI-1268 b, TOI-1683 b, TOI-2018 b, MASCARA-2 b, WASP-189 b, TOI-2046 b, TOI-1431 b, and HAT-P-57 b. We report two new highresolution spectroscopy He i detections for TOI-1268 b and TOI-2018 b, and a H α detection for TOI-1136 d. Furthermore, we detect hints of He i for HD 63433 b, and H α for HD 73583 b and c, which need to be confirmed. The aim of the Measuring Out-flows in Planets orbiting Young Stars (MOPYS) project is to understand the evaporating phenomena and test their predictions from the current observations. We compiled a list of 70 exoplanets with He i and / or H α observations, from this work and the literature, and we considered the He i and H α results as proxy for atmospheric escape. Our principal results are that 0.1-1Gyr planets do not exhibit more He i or H α detections than older planets, and evaporation signals are more frequent for planets orbiting ∼ 1-3 Gyr stars. We provide new constraints to the cosmic shoreline, the empirical division between rocky planets and planets with atmosphere, by using the evaporation detections and we explore the capabilities of a new dimensionless parameter, R He / R Hill, to explain the He i triplet detections. Furthermore, we present a statistically significant upper boundary for the He i triplet detections in the T eq versus ρ p parameter space. Planets located above that boundary are unlikely to show He i absorption signals.  \nKey words. techniques: photometric - techniques: spectroscopic - planets and satellites: atmospheres - planets and satellites: gaseous planets planets and satellites: physical evolution', '1. Introduction': "NOTE: This accepted version of the manuscript has a di ff erent numbering of figures, tables and appendices than the published version. Please refer to the published version for proper referencing.  \nAlong with their stars, planets evolve and change over time. During their early stages of formation, in particular, they undergo severe changes in their physical and orbital properties due to internal and external forces (Baruteau et al. 2016). Exoplanets form embedded in the protoplanetary disc, from where they can accrete H and He, which are the major constituents of their gaseous envelopes (Guenther et al. 2023). Following the disc dispersion in ∼ 1-10 Myr (Haisch et al. 2001; Hernández et al. 2007; Baruteau et al. 2016), di ff erent phenomena occur favouring the mass-loss processes in the planetary atmospheres: i) the gas disc dissipates leaving the exoplanet and its H / He-rich atmosphere exposed to the direct radiation of the host star (Fedele et al. 2010; Barenfeld et al. 2016; Dawson & Johnson 2018); ii) stars are more active and have their highest levels of X-ray and extreme ultraviolet (XUV, 1-912 Å) energy irradiation until ∼ 100 Myr (Jackson et al. 2012); iii) exoplanets are undergoing  \ncontraction and cooling processes until ∼ 1 Gyr (Ginzburg et al. 2016, 2018; Gupta & Schlichting 2020); iv) initial gas envelopes tend to be inflated by internal and external heat, resulting in more extended atmospheres (Sanz-Forcada et al. 2011; Owen & Wu 2016, and references therein); and v) the XUV radiation supports the population of metastable He i (Sanz-Forcada & Dupree 2008). Therefore, at this stage of the system's evolution the conditions are optimal for detecting the young primordial extended atmospheres.  \nPhoto-evaporation processes are stronger when they are driven by X-ray radiation, which is more significant during the saturated luminosity state (from ≲ 100 Myr until ∼ 1 Gyr for late spectral types; Jackson et al. 2012). After this timescale, photoevaporation is mainly driven, progressively, by the extreme ultraviolet (EUV, 100-912 Å) flux, and its e ff ects on the exoplanet atmosphere are less significant (Sanz-Forcada et al. 2011; Owen & Jackson 2012; Owen & Wu 2013). Furthermore, atmospheric escape can also be driven by the heat released from the cooling core of the planet. This mechanism, called core-powered evaporation, plays an important role, and this process can act approximately on gigayear timescales (Ginzburg et al. 2016, 2018; Gupta & Schlichting 2020).  \nBoth escape mechanisms have some overlap in time and may act simultaneously in shaping the final atmosphere of mature exoplanets. In this context, the study of planets at the early stages of evolution is crucial for a better comprehension of the di ff erent processes, and to prove theory predictions, such as planet formation and migration, giant planets' gas accretion, or escape of the primary atmospheres of rocky-core planets with gas envelopes (e.g. Baruteau et al. 2016; Owen & Wu 2017; Dawson & Johnson 2018). In particular, core-powered and photo-evaporation processes have been suggested to explain the formation of the radius gap in the small (1-4 R ⊕ ) planets radius distribution (e.g. Fulton et al. 2017; Fulton & Petigura 2018; Van Eylen et al. 2018, 2021), although formation and migration mechanisms have also been proposed to explain the valley (Luque & Pallé 2022).  \nExoplanets can also experience atmospheric mass loss during their later lifetime because photo-evaporation and corepowered mechanisms can act at longer timescales than ∼ 0.11 Gyr, but at reduced total mass-loss rates (Sanz-Forcada et al. 2011; Owen & Jackson 2012). According to stellar and planetary evolution models and assuming that high mass-loss rates are easier to detect, the number of evaporation detections should be higher in the < 100Myr period, where the X-ray photoevaporation and core-powered mechanisms act together (SanzForcada et al. 2011; Owen & Wu 2013; Ginzburg et al. 2016). After ∼ 1 Gyr, the number of detections should decrease as the core-powered atmospheric escape is weakening during the first few gigayears (Ginzburg et al. 2016, 2018). As a rough approximation, we should expect a plateau of atmospheric massloss detections after the approximately gigayear timescale. At the time photo-evaporation is mainly driven by the EUV flux (Sanz-Forcada et al. 2011), the planet properties are already fixed. For gas giants and Jupiter-mass planets, the impact of this late-evaporation on the planet mass is close to null (Owen & Wu 2013), while for low-mass planets only small changes in their masses are expected (Lopez et al. 2012; Owen & Jackson 2012; Owen & Wu 2013). Therefore, the detection of hydrodynamic escape in old planets ( ≳ 1 Gyr) may not imply a significant change in their total mass and it may have no evolutionary consequences. Hydrodynamic escape by Roche lobe overflow can be considered an exception (see e.g. Koskinen et al. 2022). The middle atmosphere of extremely close-in planets (with orbital periods of ≲ 1 d) extends to the Roche lobe due to the stellar gravitational tide, producing a very strong atmospheric massloss rate. The Roche lobe mechanism can be assumed stellar-age independent, and may not introduce a detectable change in the number of detections across stellar age.  \nThe study of atmospheric escape in exoplanet atmospheres is mainly performed with the observations of three evaporation tracers: the H Lymanα line (Ly α ) at 1216 Å at ultraviolet (UV) wavelengths, the H Balmerα line (H α ) at 6564 Å at optical wavelengths, and the He i triplet at 10833 Å at near-infrared (NIR) wavelengths. All the wavelengths in this work are referenced in a vacuum.  \nThe Ly α was the first line used for probing evaporating atmospheres, when Vidal-Madjar et al. (2003) detected its absorption in the upper atmosphere of HD 209458 b. However, the Ly α line has some limitations. First, UV observations cannot be carried out from ground-based facilities, relegating its study to space observations with the Hubble Space Telescope (HST) instrument Space Telescope Imaging Spectrograph (STIS). Second, Ly α is strongly a ff ected by the interstellar medium extinction, limiting its observation to only the closest stars, due to their large relative motions. Finally, the H in the upper layers of Earth's atmosphere  \n(geocorona) have Ly α emission. In practice, the core of the Ly α line cannot be observed, and only the broad wings of the line are accessible. This has led to a very limited number of planetary atmospheres being accessible to Ly α observations.  \nThe H α line was used by Yan & Henning (2018) to prove the evaporation of KELT-9 b's atmosphere. The core and wings of this optical line can be observed from space, but also from ground-based facilities, allowing line profile characterisation when high-resolution spectrographs are used. However, the H α line has some caveats as well: relatively low mass-loss rates can produce large Ly α absorption signals, while remaining undetected in H α (e.g. GJ 436 b, Kulow et al. 2014; Ehrenreich et al. 2015; Cauley et al. 2017); stellar lines in the visible are more sensitive to stellar activity than those in the NIR (e.g. H Paschen lines and He i triplet), making their analysis very challenging in the presence of such activity (Fuhrmeister et al. 2020; Palle et al. 2020b; Howard et al. 2023).  \nThe He i triplet was proposed by Seager & Sasselov (2000), and later modelled by Oklopˇci'c & Hirata (2018), as an alternative to search for evidence of atmospheric escape in the NIR. The first detections came almost simultaneously from groundand space-based observations (Spake et al. 2018; Nortmann et al. 2018; Allart et al. 2018). Moreover, ground-based highresolution spectroscopy observations allowed several physical parameters to be retrieved from line profile fitting (Lampón et al. 2020, 2021b,a, 2023). Since stellar lines in the NIR are less sensitive to stellar activity (but not exempted from it; e.g. Spake et al. 2018; Salz et al. 2018) than the optical lines, the He i triplet is a good tracer to explore young exoplanet atmospheres.  \nUnfortunately, the He i triplet has some disadvantages too. First, there are telluric absorption and emission lines surrounding the NIR triplet wavelengths that can overlap with the planet signal depending on the observing epoch. Second, the He i ionisation wavelength cuto ff ( λ ≤ 504 Å) is lower than that of H i ( λ ≤ 912 Å). Finally, the lifetime of the helium metastable state is ∼ 2.2 h (Drake 1971), and a strong and constant XUV irradiation is needed to maintain the helium metastable state population detectable (Sanz-Forcada & Dupree 2008).  \nThis work presents the first results of the Measuring Outflows in Planets orbiting Young Stars (MOPYS) survey. The MOPYS project's aim is to provide observational constraints to the timescales of atmospheric evolution and mass loss processes, and test the predictions of planetary evolution models by comparing them to the evaporation tracers observations. In particular, we are interested in testing the mass-loss processes in planet formation at di ff erent timescales. Assuming that young planets undergo higher mass-loss rates than older ones, and higher mass-loss rates are, in general, easier to detect, we search for signs of evaporation using the He i triplet and / or H α line observations as proxy, focusing on transmission spectroscopy observations of ≲ 1 Gyr exoplanets. Moreover, we study the detection rate between young and old planet populations in our sample, and compare the observed He i triplet signals with predicted photo-evaporation mass-loss rates. We scheduled the observations such that telluric contamination of the He i triplet is minimised, taking advantage of a favourable barycentric velocity of the Earth. The majority of the high-resolution spectrographs that we used allowed us to observe the He i triplet and H α simultaneously: CARMENES (Quirrenbach et al. 2014, 2020), HARPS-N + GIANO-B(GIARPS; Claudi et al. 2016), and HARPS + NIRPS (Mayor et al. 2003; Wildi et al. 2022). Although H α might be more a ff ected by the probable stellar activity, the simultaneous observation of the two lines increases  \nthe capability to detect ongoing evaporation. For consistency, we complemented the database of He i triplet and H α observations with some Ly α observations from the literature when possible, although this work does not focus on this UV line.  \nThis manuscript is organised as follows. We describe the observational datasets and the spectroscopic analysis methodology in Section 2. The atmospheric results for each planet are detailed in Section 3, while Section 4 introduces further H and / or He i atmospheric results from the literature. In Sect. 5 we detail our criteria to consider atmospheric detections. In Section 6 we discuss our main findings regarding atmospheric evaporation and the He i triplet. The conclusions of this work can be found in Section 7.", '2.1. Spectroscopic observations': 'In this work we analysed 16 transits observed with the Calar Alto high-Resolution search for M dwarfs with Exoearths with Near-infrared and optical Échelle Spectrographs (CARMENES, Quirrenbach et al. 2014, 2020) spectrograph located at the Calar Alto Observatory, Almería, Spain. CARMENES has two spectral channels: the optical channel (VIS), which covers the wavelength range from 0.52-0.96 µ m with a resolving power of R = 94 600, and the near-infrared channel (NIR), which covers 0.96-1.71 µ m with a resolving power of R = 80 400. The targets were observed with both channels simultaneously.  \nFibre A was used to observe the targeting star, while fibre B was placed on blank sky in order to monitor the sky emission lines (fibres A and B are separated by 88 arcsecs in the east-west direction). The observations were reduced using the CARMENESpipeline caracal (Caballero et al. 2016), and both fibres were extracted with the flat-optimised extraction algorithm (Zechmeister et al. 2014).  \nWe also analysed 11 transits observed with the High Accuracy Radial velocity Planet Searcher for the Northern hemisphere (HARPS-N, Cosentino et al. 2012) and / or GIANO-B (Oliva et al. 2012) spectrographs mounted on the 3.6m Telescopio Nazionale Galileo (TNG) at Roque de los Muchachos Observatory, La Palma, Spain. HARPS-N is an optical spectrograph which covers the wavelength range from 0.383-0.693 µ m with a resolving power of R = 115 000. HARPS-N spectra were extracted using the standard Data Reduction Software (DRS) pipeline (Cosentino et al. 2014). GIANO-B (Oliva et al. 2012) is a near-infrared spectrograph which covers the wavelength range from 0.95-2.45 µ mwith a resolving power of R≃ 50 000. When possible, the observations were done in GIARPS mode (Claudi et al. 2016), which allows the simultaneous use of HARPS-N and GIANO-B spectrographs.  \nGIANO-B observations were carried out with the nodding acquisition mode, where the object is observed at two di ff erent predefined positions on the slit (A and B) following an ABAB pattern (Claudi et al. 2016). The nodding technique enable to monitor the sky with the slit position that is not on the object, and then e ffi ciently subtract the thermal background and telluric emission lines. The GIANO-B spectra were wavelength calibrated, and extracted using the GOFIO pipeline (Rainer et al. 2018).  \nTable 1 shows the observing log of the planetary transits analysed in this work, indicating the instrument used in each case. We computed the central time of transit ( T c) and its uncertainty ( σ T c ) with the Transit and Ephemeris Service tool from  \nNASA Exoplanet Archive 1 and the parameters from Table 2. The core of the MOPYS observations were conducted as part of the 21B-3.5-004, 22A-3.5-007, 22B-3.5-009, 23A-3.5-009, and 23B-3.5-002 observing programmes (PI J. Orell-Miquel) with CARMENES, and CAT21B\\_61, CAT22A\\_9, CAT22B\\_43, and CAT23A\\_100 observing programmes (PI J. Orell-Miquel) at TNG. K2-100 b, V1298 Tau c, HD 235088 b, and TOI-2048 b were observed using GTO time by the CARMENES consortium.  \nWe processed the VIS and NIR CARMENES, and HARPSN observations with serval 2 (Zechmeister et al. 2018), which derives the relative radial velocities (RVs) and several activity indicators: the chromatic radial velocity index (CRX), the differential line width (dLW), and the H α , Na i D1 and D2 and Ca ii IRT line indices. Moreover, for the HARPS-N datasets we also used the YABI tool, 3 an online version of the HARPS-N DRS pipeline, to derive absolute RVs and spectral activity indicators: cross-correlation function (CCF) full width at half maximum(FWHM),CCFconstrast (CTR), bisector (BIS), and MontWilson S-index. For the targets only observed with GIANO-B (because the GIARPS mode was not possible), we inspected the H Paschen lines in the NIR as stellar activity indicators. We constructed the stellar light curve of the H Paschen β (Paβ , 12821.6 Å), H Paschen γ (Paγ , 10941.1 Å), and H Paschen δ (Paδ , 10052.1 Å) lines. Abrupt and / or strong variations in the time evolution of the activity indices or lines may indicate stellar activity or flares during the transit, and could compromise or challenge the detection of planetary signals, in particular absorption lines in the visible part of the spectrum (Palle et al. 2020b; Orell-Miquel et al. 2023).  \nRV measurements during a transit can be used to look for the Rossiter-McLaughlin (RM; Rossiter 1924; McLaughlin 1924) e ff ect. During the crossing of a planet in front of its host star, the planet blocks di ff erent parts of the stellar disc. That produces an RV anomaly during the transit known as the RM e ff ect. The detection of the RM e ff ect in RV time series taken during a planetary transit enables the confirmation of the presence of the transiting planet and also helps to determine the orbital configuration and architecture of the planetary system. This technique has been successfully applied to young planets unveiling the architecture of its planetary systems, such as AU Mic b (Palle et al. 2020b) or DS Tuc A b (Benatti et al. 2021). Because we used the CARMENES instrumental configuration where fibre B points at the sky, there were no simultaneous Fabry-Pérot calibrations during the observations. Without them, the RV time evolution is dominated by the instrument drift during the night.', '2.2. Photometric observations': 'We refined the ephemerides and planetary properties of TOI2048 b, HD 73583 b & c, K2-77 b, TOI-1807 b, TOI-1683 b, TOI-2076 b, TOI-2018 b, and TOI-1268 b using photometric data.  \nFrom the Transiting Exoplanet Survey Satellite (TESS; Ricker et al. 2015), we analysed the 2-minute cadence TESS simple aperture photometry (SAP; Morris et al. 2017) using juliet 4 (Espinoza et al. 2019). This python library is based on other public packages for transit light curve ( batman , Kreidberg 2015), and for Gaussian process (GP; celerite , Foreman-  \nTable 1: Observing log of the transits analysed in this work. Columns from left to right: name of the planet, spectrograph used, date of starting night of the observations, starting and ending time of observations, central time of the scheduled transit T c, 1 σ uncertainty of T c, median exposure time of the observations, number of spectra fully taken between the first ( T 1) and fourth ( T 4) contacts ( N T14) compared to the total number of spectra taken during the observations ( N obs), median signal-to-noise ratio (S / N) value at the spectral orders of H α and / or He i lines, respectively.  \nNotes. ( a ) No planetary transit was observed.  \nMackey et al. 2017) modelling and uses nested sampling algorithms ( dynesty , Speagle 2020; MultiNest , Feroz et al. 2009; Buchner et al. 2014) to explore all the parameter space. In the fitting procedure, we adopted a quadratic limb-darkening law with the ( q 1, q 2) parameterisation introduced by Kipping (2013), and we considered the uninformative sample ( r 1, r 2) parameterisation introduced in Espinoza (2018) to explore the impact parameter of the orbit ( b ) and the planet-to-star radius ratio ( p = R p / R ⋆ ) values. We modelled the photometric variability of the young host stars adding to fit the celerite GP exponential or celerite GP quasi-periodic kernels. We followed the same procedure applied to analyse the HD 235088 system in OrellMiquel et al. (2023).  \nOne transit of TOI-1268 b and TOI-2018 b each were observed by the multi-colour imager MuSCAT2 (Narita et al. 2019) mounted at the 1.52-m Telescopio Carlos Sánchez (TCS) in the Teide Observatory in Tenerife, Spain. MuSCAT2 obtained simultaneous photometric data of the transits in four bands (Sloan g , r , i , zs ). The exposure times for each band were optimised for each night, CCD and target to avoid the saturation of the target and comparison stars in the field. Standard data reduction, aperture photometry, and transit modelling including systematic e ff ects was performed by the MuSCAT2 custom pipeline (Parviainen et al. 2019).  \nOne partial transit of TOI-2076 b was observed on the night of 14 May 2023 with the Sinistro imager mounted on one of the two 1-m telescopes (Dome B) operated by Las Cumbres Ob-  \nservatory (LCO; Brown et al. 2013) at the Teide Observatory in Tenerife, Spain. The observations were performed through the rp-band filter, in the full-frame mode, and with an exposure time of 10 s. To avoid detector saturation, the telescope was defocused such that the FWHM of the stellar point spread function was 7"8". The observation started at 22:00 UT, about 1.6 hours before the expected ingress time, and halted at 01:03 UT, in the middle of the expected transit, due to a technical problem. The obtained images were processed by the BANZAI pipeline (McCully et al. 2018) for dark and flat corrections. The light curve of TOI2076 was then extracted by aperture photometry using a custom pipeline (Fukui et al. 2011), and the transit modelling was performed with juliet .  \nMore details of each transit analysis can be found in their corresponding sections (TOI-1268 b: Sect. 3.15, TOI-2018 b: Sect. 3.20, and TOI-2076 b: Sect. 3.16).', '2.3. Telluric correction': "The main objective of these observations was the analysis of the He i triplet at ∼ 10833 Å. However, the He i lines are surrounded by emission and absorption telluric lines. There is an H2O absorption line at 10835 Å and there are four emission lines of hydroxyl (OH) at 10832.1 Å, 10832.4 Å, 10834.2 Å, 10834.3 Å (Oliva et al. 2015), although the last two are detected in the spectra as a single peak. Due to the Earth's orbital motion the relative positions of the planet and telluric spectral lines change with the epoch. Thus, we only scheduled the observations during transiting epochs that minimise overlapping and mitigate the contamination over the He i triplet (see Orell-Miquel et al. 2022; Spake et al. 2022 for further details).  \nWhen analysing the transits presented in this work, we encountered another source of contamination of the He i triplet. We detected an emission line at ∼ 10833 Å in few transits where we needed to extend the observations into the astronomical twilight. This emission line only appears in the last few spectra and its strength increases towards the end of the night. We presume it is scattered sunlight from the incoming dawn.  \nBecause each spectrograph has its own particularities, we handled with the telluric emission and absorption contamination di ff erently for each of the instruments, see next subsections.", '2.3.1. CARMENES telluric correction': 'We corrected the CARMENES VIS and NIR spectra from telluric absorptions following the approach described in Nagel et al. (2023) with the molecfit package in version 1.5.9 (Smette et al. 2015; Kausch et al. 2015). Then, we corrected the He i region of the CARMENES NIR data from telluric OH emission following the methodology described in previous He i studies (e.g. Nortmann et al. 2018; Palle et al. 2020a; Czesla et al. 2022), in particular Orell-Miquel et al. (2022, 2023).  \nWe used fibre B to generate a synthetic emission model fitting simultaneously the three main OH peaks with three independent Gaussian profiles. The amplitude, central position, and standard deviation of the Gaussian profiles were set free, and we only introduced an initial guess for their central positions. Prior to applying the emission model to fibre A, we accounted for the di ff erent e ffi ciency of the two injection fibres. For each dataset, we computed the scaling factor between fibres comparing the strongest OH peak from the co-added spectra of each fibre. We assumed the scaling factor to be constant during the night. Finally, for each pair of target and sky spectra, we divided  \nthe science spectra by its particular OH emission model, multiplied by the nightly scaling factor. When we detected the solar He i emission line in fibre B, we added a fourth Gaussian profile in the fitting procedure.', '2.3.2. HARPS-N telluric correction': 'We corrected the HARPS-N spectra from telluric absorptions with molecfit (v4.2) via its implementation in the SLOPpy (Spectral Lines Of Planets with python, Sicilia et al. 2022) package. We only used SLOPpy for the purpose of running molecfit easily on HARPS-N data, and correct those spectra from telluric absorptions.  \nAlthough the wavelength solution from the original HARPSN data is based on wavelengths in air, the molecfit correction provides a wavelength solution in the vacuum. Because CARMENESand GIANO-B wavelengths are in the vacuum, we also used the vacuum wavelengths for the HARPS-N data for consistency.', '2.3.3. GIANO-B telluric correction': 'Due to our planning of the observations, the H2O absorption line is always far from the He i triplet. Thus, we decided to not correct the GIANO-B spectra from telluric absorption. If a particular residual map or transmission spectrum shows strong telluric residuals, we simply masked that spectral region.  \nThe sky emission lines are corrected by the GOFIO pipeline, taking advantage of the nodding technique. This procedure is very e ffi cient in removing emission lines that have similar strength in consecutive exposures. However, the He i emission increases too quickly between exposures for the ABAB nodding procedure to provide an accurate correction for the emission contamination. We inspected the individual spectra and masked the a ff ected wavelength region of particular spectra.', '2.4. Transmission spectrum analysis': "We analysed the spectroscopic observations via the wellestablished transmission spectroscopy technique (e.g. Wyttenbach et al. 2015; Casasayas-Barris et al. 2017) successfully applied in Orell-Miquel et al. (2022, 2023). The planetary and stellar parameters needed to compute the transmission spectra for each planet are shown in Table 2.  \nThe telluric corrected spectra are normalised by a first degree polynomial. We used a blue and red region near the line of interest and free of tellurics to fit the polynomial. The spectra are shifted into the stellar rest frame, accounting for the Earth's barycentric movement, the stellar systemic velocity ( γ ), and the stellar Doppler shift induced by the planet. Then, we created a high signal-to-noise ratio (S / N) stellar spectrum (Master-Out; MO) by calculating the mean of all the spectra taken completely during out-of-transit, and we divided all the spectra by the MO. With this step, the stellar spectrum is removed from the spectroscopic time series. Next, we moved the spectra to the planetary rest frame using the formula  \nv p( t ) = K p sin (2 πϕ ( t )) , (1)  \nwhere K p is the radial velocity semi-amplitude that the star induces to the planet, and ϕ is the orbital phase. We neglected orbital eccentricity because the values provided for the planets in the literature are usually poorly constrained and the impact  \nTable 2: Transit and system parameters used to compute the transmission spectra for each planet analysed in this work.  \nReferences. K2-100b: (1) Stefansson et al. (2018), (2) Barragán et al. (2019). MASCARA-2 b: (1) Lund et al. (2017) , (2) Talens et al. (2018), (3) Casasayas-Barris et al. (2019). V1298 Tau c: (1) J. Livingston, priv. comm., (2) David et al. (2019), (3) Suárez Mascareño et al. (2021), (4) Gaidos et al. (2022). TOI-1431 b: (1) Addison et al. (2021), (2) Stangret et al. (2021). TOI-2048 b: (1) This work (App. C), (2) Newton et al. (2022) HD 63433 b: (1) Mann et al. (2020), (2) Zhang et al. (2022c), (3) Mallorquín et al. (2023) HD 63433 c: (1) Mann et al. (2020), (2) Zhang et al. (2022c), (3) Mallorquín et al. (2023) HD 73583 b: (1) Barragán et al. (2022), (2) El Mufti et al. (2021). HD 73583 c: (1) Barragán et al. (2022), (2) El Mufti et al. (2021). HD235088b: (1) Orell-Miquel et al. (2023), (2) Gaia DR2 (Soubiran et al. 2018). K2-77 b: (1) This work (App. F), (2) Gaidos et al. (2017). TOI2046 b: (1) Kabáth et al. (2022), (2) Computed using this table parameters. TOI-1807 b: (1) This work (App. G), (2) Nardiello et al. (2022), (3) Gaia DR2 (Soubiran et al. 2018). TOI-1136 d: (1) Dai et al. (2023a), (2) Gaia DR2 (Soubiran et al. 2018). TOI-1268 b: (1) Šubjak et al. (2022), (2) This work (Sect. 3.15), (3) Gaia DR2 (Soubiran et al. 2018). TOI-2076 b: (1) Zhang et al. (2023b), (2) This work (Sect. 3.16), (3) Osborn et al. (2022), (4) Gaia DR2 (Soubiran et al. 2018). TOI-1683 b: (1) This work (Sect. K) , (2) Gaia DR2 (Soubiran et al. 2018), (3) Zhang et al. (2023b). WASP-189 b: (1) Lendl et al. (2020) , (2) Anderson et al. (2018). HAT-P-57 b: (1) Hartman et al. (2015). TOI-2018 b: (1) Dai et al. (2023b), (2) This work (Sect. 3.20), (3) Gaia DR2 (Soubiran et al. 2018).  \nNotes. ( a ) T 0 given in BJD -2 457 000, except for HAT-P-57 b that is BJD -2 450 000. ( b ) T 14 is the total transit duration between the first ( T 1) and fourth ( T 4) contacts, and T 12 is the duration of the ingress-egress. ( c ) Calculated from K p = 2 π a p P -1 sin i p using the parameters in this table, when there is no reference.  \nfor all but one case is of the same order as the uncertainties derived for the signal shifts. The exception to this is TOI-1268 b for which we used the radvel (Fulton et al. 2018) package to account for the eccentricity shift during the transit.  \nTo compute the transmission spectrum (TS), we only consider those spectra taken entirely between the first ( T 1) and fourth ( T 4) contacts. We averaged them using the inverse of the squared propagated errors (1 / σ 2 ) as weights. If the residual map contains strong telluric residuals and / or stellar line variability near the inspected line, we mask those a ff ected regions to compute the TS. However, each night is di ff erent, and we explain further details of their particularies in the data analysis section (Sect. 3), where needed.  \nWhen the TS shows a clear absorption at the expected position, we fitted the feature with a Gaussian profile using the MultiNest algorithm (Feroz et al. 2009) via its python implementation PyMultinest (Buchner et al. 2014). Otherwise, when the residual map or the transmission spectrum do not show evidences of absorption from the planetary atmosphere, we derived conservative limits for the planetary absorption. We estimated a 3 σ upper limit absorption peak as three times the root-meansquared (RMS) value of flat spectral region in the continuum of the spectrum close to the line(s) of interest. The exact spectral range might be di ff erent in each case to avoid including strong  \nresiduals or very scattered regions. The equivalent width (EW) of a line is a more appropriate measure because it is independent from the instrumental resolution, but it is di ffi cult to compute when there is no line to measure. We give the upper limits in terms of absorption peak and we translate them into EWs (or from EW to absorption peak) as it is explained in App. M.  \nTo analyse the data from GIANO-B, we followed some of the recommendations from Guilluy et al. (2020), Fossati et al. (2022), and Guilluy et al. (2023). We considered the science spectra taken at positions A and B as di ff erent instruments, and we combined them after computing both TS. GIANO-B is known to be a ff ected by a fringing pattern that a ff ects high S / N spectra (Guilluy et al. 2020). However, the lower S / N of our spectra, due to its maximum exposure time of 600 s and the faintness of our targets, prevented us to detect hints of the fringing pattern. The MO spectra obtained from GIANO-B did not show any clear sinusoidal pattern, as seen in Guilluy et al. (2020, Fig. 2). Thus, we decided to not correct the fringing pattern, as it was done in Fossati et al. (2022). Lastly, the wavelength solution of GIANO-B may not be stable during the night. Guilluy et al. (2020) found a typical instrument drift of half of the GIANO-B pixel element during four hours of observations. That could be an issue for the cross-correlation technique, where excellent precision in the wavelength solution is needed (Brogi  \net al. 2018). However, the transmission spectra technique does not require such extreme precision, and hence we did not apply any correction to our GIANO-B observations, which have lower S / N than in Guilluy et al. (2020).  \nAnother aspects to take into account with transit observations are the Rossiter-McLaughlin (RM) and center-to-limb-variation (CLV) e ff ects, which can interfere or mimic single line absorptions (Casasayas-Barris et al. 2021b). Because these e ff ects scale with R p / R ⋆ , they have larger impact on gas giant transmission observations. Previous He i triplet line works proved that they have negligible contributions on the He i results (e.g. Nortmann et al. 2018; Guilluy et al. 2023). Here, we computed the RM and CLV models for HD 63433 b ( R p = 2.1 R ⊕ ) and TOI-1136 d ( R p = 4.6 R ⊕ ) whose contribution to the H α and He i triplet signals is negligible. These analyses and results are presented in Sect. 3. The models of RM and CLV e ff ects were created with the Turbospectum2019 (Plez 2012), using the Kurucz ATLAS9 (Castelli & Kurucz 2003) and VALD3 (Ryabchikova et al. 2015) line lists. The stellar models were created for 21 limb-darkening angles assuming local thermodynamical equilibrium (LTE). In the next step, taking into account the system parameters, we calculated the stellar models for orbital phases from -0.05 to 0.05, remembering that the planet covers di ff erent stellar regions with di ff erent limb-darkening during its transit. Next, all models were divided by the out-of-transit spectrum, creating the RM and CLV e ff ect model for each orbital phase and wavelength. However, it is necessary to first detect the RM in the RV time series to compute the models, and the e ff ect is not detected for all our targets.", '3. Transmission spectroscopy results': 'In this section we present the results and the details of the residuals maps and TS calculations planet by planet. We present an example of a residuals map and TS figure in Figure 1, and the figures for the rest of transits without detections are shown in Appendix A. Furthermore, the time evolution of the activity indicators derived from the analysed transits can be found in Appendix A as well.', '3.1. K2-100 b': 'The resulting residual maps and TS centred in the spectral regions of the H α and He i triplet are shown in Figure 1. The time evolution of all the activity indices is mainly flat, and there is no evidence of strong stellar variability during the transit.  \nThe He i triplet TS shows a flat spectrum, while the H α TS is flat with some emission features from stellar variations in the line core during the transit. We found no significant absorption in either of the lines, and we placed a 3 σ upper limit to the excess absorptions of 1.4 % and 1.3 % for H α and He i , respectively.  \nK2-100b was already observed with the InfraRed Doppler (IRD) spectrograph on the 8.2-m Subaru telescope ( R = 70 000; Kotani et al. 2018). Gaidos et al. (2020a) derived an EW 99% confidence detection upper limit of 5.7 mÅ. From Fig. 8 therein, we estimated an absorption depth of ∼ 1.2%, which is consistent with the upper limit derived from the CARMENES observations.', '3.2. MASCARA-2b': 'Casasayas-Barris et al. (2019) already detected H α in MASCARA-2b using CARMENES and HARPS-N observations. The H α absorption measured with CARMENES VIS data  \nwas -0.85 ± 0.03 %, and while combining 3 nights with HARPSN was -0.68 ± 0.06 %.  \nHere, we inspected the CARMENES NIR spectra, which were not analysed in Casasayas-Barris et al. (2019), looking for He i excess absorption. Figure A.1 presents the residual map and TS around the He i triplet. Because some telluric OH variability was still remaining in the residual map, we masked the a ff ected regions. Despite the structure in the final TS, due to the masked regions, we found no significant planetary absorption. We computed a 3 σ upper limit to the He i excess absorption of 0.5 %.', '3.3. V1298 Tau c': 'V1298Tau planetary system is known to have transit time variations (TTVs) that can complicate the analysis of high-resolution spectroscopy observations (David et al. 2019). According to non-linear ephemerids for the V1298 Tau system (J. Livingston, priv. comm.), a nearly complete transit of planet c was observed the night of 5 of January 2020. Unfortunately, we missed pretransit observations which are critical to compute a good MO and to check for stellar variability during the observations.  \nFor this transit, the serval H α index shows a clear decrease only from about mid-transit to egress, and an increase at the end of the observations. We did not include those last spectra in the computation of the MO spectrum, although they are plotted in the residual maps. The other serval products do not show signs of strong stellar activity (see Fig. A.20).  \nThe H α residual map (Fig. A.2 left) shows a clear absorption region coincident with the serval H α index decrease, suggesting the stellar origin of the signal. We computed the TS only with the non-a ff ected spectra from the first half of the transit. We fitted a -1.10 ± 0.13 % feature at the H α position, which is consistent with the 3 σ upper limit derived from the nearby spectral region (1.1 %). The priors and posterior values from the fit are listed in Table B.1, and the posterior distributions are shown in Fig. B.1.  \nThe first half of the transit in the He i NIR triplet residual map (Fig. A.2 left) shows an extended absorption region, which does not coincide in time with the H α variability detected. In the second half of the transit, there is a vertical absorption feature in the stellar rest frame that does not follow the planetary track when we shift the spectra into the planet rest frame. This feature seems to be related to telluric residuals. We computed the TS with those una ff ected spectra from the first half of the transit. The He i TS has a deep ( -3.75 ± 0.12 %) and broad (1.02 ± 0.04 Å) absorption feature, which is larger than the 3 σ upper limit (1.1 %). The priors and posterior values from the fit are listed in Table B.2, and the posterior distributions are shown in Fig. B.1.  \nTo explore the origin of the H α and He i features, we computed the TS from the una ff ected spectra (start phase), the centre of the transit (centre phase), the end of the transit (end phase), and post-transit. We compare those TS with the MO spectra in Fig. B.2. The start TS from H α and He i lines show a di ff erent behaviour from centre and end. Furthermore, the comparison between TS and MO evidences the stellar line profile variation during the observations. Therefore, we cannot confidently attribute the detected features to V1298 Tau c atmosphere. Moreover, the H α TS in Fig. B.2 show a tiny bump at ∼ 6563 Å which can be attributed to the star and its youth. The surprisingly large ∆ v ( -28.7 ± 1.1 km s -1 ) is another reason for not claiming the significant He i signal as a detection.  \nThe H α line of V1298 Tau c was previously explored by Feinstein et al. (2021) and Schlawin et al. (2021). Feinstein et al.  \nFig. 1: Residuals maps and transmission spectra around the H α line ( top panels ) and He i NIR triplet ( bottom panels ) lines for K2-100b observations with CARMENES. Left panels : Residual maps in the stellar rest frame. Time since mid-transit time ( T c) is shown on the vertical axis, wavelength is on the horizontal axis, and relative absorption is colour-coded. The dashed and dotted white horizontal lines indicate the di ff erent contacts during the transit. The dashed cyan tilted lines indicate the predicted trace of the planetary signals. The solid green vertical lines indicate the position of the OH emission telluric lines. The solid blue line indicates the position of the H2O absorption telluric line. Right panels : Planet transmission spectra (TS) in the planet rest frame. We show the original data in light grey and the data binned by 0.2 Å in black. When an absorption signal is fitted, a red line and shaded region show the best Gaussian fit model with its 1 σ uncertainties. The dotted cyan vertical lines indicate the H α ( top ) and the He i triplet ( bottom ) lines positions. All the wavelengths in this figure are referenced in a vacuum.  \n<!-- image -->  \n(2021) studied with the Gemini / GRACES spectrograph the behaviour of the H α and Ca ii infrared triplet during a transit, finding significant variations of the H α during the observations. Schlawin et al. (2021) analysed one transit looking for H α absorption with the PEPSI spectrograph, but the observations were a ff ected by stellar variability. Furthermore, Vissapragada et al. (2021) explored the He i triplet of planet c with the Habitablezone Planet Finder (HPF) spectrograph, but the in-transit spectra were a ff ected by a stellar flare. Lastly, a transit of planet c was observed with narrow-band helium filter photometry, but Vissapragada et al. (2021) did not detect the transit during the observations.  \nUnfortunately, our observations join the previous unsuccessful attempts to analyse the atmosphere of V1298 Tau c. Due to the many complications that a ff ected the observations, we adopted the absorption peaks from the H α and He i features as upper limits of any possible planetary absorption. Our upper limits to H α and He i excess absorption are 1.1 % and 3.7 %, respectively.', '3.4. TOI-1431 b': 'Stangret et al. (2021) analysed two HARPS-N and one EXPRES transits, finding no signs of atomic or molecular absorption. A tentative detection of H α appeared only during the first HARPSN dataset (0.33 ± 0.07 %), but it was considered as stellar variability because it was not reproduced in the other two visits. We consider that absorption feature as an upper limit for the purpose of this work.  \nHere, we looked for He i excess absorption in the GIANO-B spectra from the first night of Stangret et al. (2021). The residual map and TS are shown in Fig. A.3. The TS is mainly flat and we placed a 3 σ upper limit to the He i excess absorption of 0.4 %.', '3.5. TOI-2048 b': 'Newton et al. (2022) noted that ground-based follow-up observations found a tentative transit detection ∼ 20 min later than expected. We analysed the new available TESS data to derive updated ephemerides. The details of the photometric fit are explained in Appendix C. We refined TOI-2048 b ephemerides and planet properties (Table C.1), and confirmed the detected delay from ground-based observations. Furthermore, because TOI2048 b does not have mass measurements, we estimated its probable mass from its radius (Wolfgang et al. 2016) to predict its semi-amplitude K ⋆ . This parameter is required for a proper measurement of the TS. We forecasted a mass of 9.4 ± 1.0 M ⊕ and a K ⋆ of 2.8 ± 0.3 m s -1 . The details of all these calculations are explained in App. C as well.  \nThe time evolution of all the activity indicators is mainly flat except for the last four H α index points. We excluded those spectra from the MO combination, although they are plotted in the residual maps. Furthermore, we masked a region near the He i triplet because some residuals still remained after correcting the OHemission. We obtained similar results with or without applying the OH emission correction.  \nH α shows a narrow absorption feature (see Fig. A.4), but we consider it too narrow to have a planetary origin (it is only one binned point). Moreover, there are emission signals around the  \nabsorption signal in the residual map, and the absorption peak is of the order of the 3 σ upper limit of 1.5 %, and is di ffi cult to distinguish from the continuum in the final TS. We found no significant absorption features around He i triplet (Fig. A.4). We placed a 3 σ upper limit to the He i excess absorption of ∼ 1 %.', '3.6. HD 63433 b': 'The night on 1 November 2021 was cloudy until ∼ 01:40 UT, and we could only cover the transit partially, without pre-transit information. The observations are a bit noisy because the exposures were set to 1 minute and we missed pre-transit observations due to bad weather conditions. Then, we observed a complete transit one year later, on 27 November 2022. The time evolution of all the activity indicators is mainly flat, without evidence of stellar activity during the observations in both nights. Although the RM e ff ect was detected by Mann et al. (2020), the RM and CLV contribution to the H α and He i triplet transmission spectra is well within the error bars of the TS (see Fig. D.1), so we decided to not correct the data from these marginal e ff ects.  \nThe H α transmission spectrum from the first partial transit displays an absorption feature of -1.21 + 0 . 22 -0 . 24 % (Fig. A.5). Table D.1 shows the priors and posterior values from the fit and Fig. D.2 shows the posterior distribution. However, we did not recover that H α absorption in our second visit, which covered the entire transit. We consider the H α feature does not have a planetary origin. Due to the di ff erent quality around the H α from both nights, we did not combine the two datasets. The residual map and TS for the H α from the second transit is shown in Fig. A.5. Weplaced a 3 σ upper limit to the H α excess absorption of 0.4 %.  \nThe residual maps and TS for the He i triplet from both nights separately are shown in Fig. A.5. The He i triplet was completely uncontaminated from tellurics, and both transmission spectra show some residuals from the stellar Si i line at 10830 Å. Due to the di ff erent quality of the nights, we did not combine them and focused on the second event which is less noisy. The TS presents a ∼ 0.25 % excess absorption feature just at the expected position, but there is noise structure of similar amplitude. Finally, we adopted a conservative 3 σ upper limit of 0.4 %.  \nOur results on the He i triplet are consistent with the previous upper limit of ∼ 0.5 % reported by Zhang et al. (2022c). Their partial transit also showed a ∼ 0.2% excess absorption, but they found a correlation between the stellar He i and H Paschen γ lines. We also compared the H Paschen lines with the He i line, but we could not explain the He i triplet feature as stellar line variability. Further high-S / N observations of HD 63433 b are required to confirm the interesting feature of ∼ 0.2% around the He i triplet found in this work and Zhang et al. (2022c).  \nAlthough the study of Ly α observations is not the focus of this work, it is worth mentioning Zhang et al. (2022c) presented HST / STIS observations of HD 63433 b around the Ly α line, but no absorption was found for this planet.', '3.7. HD 63433 c': "The time evolution of the serval activity indicators is flat, except the H α line index that displays a narrow peak close to midtransit, which is also visible in the H α residual map (Fig. A.6). Therefore, we excluded those spectra to compute the H α TS, which is mainly flat. The H α TS displays a too narrow absorption feature that originates only from the first half of the transit. Thus, we attributed this signal to a stellar origin rather than to  \nabsorption from the planet. We put a 3 σ upper limit of 0.4 % to any possible H α excess absorption from the planet's atmosphere.  \nThe He i triplet residual map and TS do not show features from the planetary atmosphere, or from stellar variability (Fig. A.6). We placed a 3 σ upper limit to the He i excess absorption of 0.4 %, consistent with the previous upper limit of ∼ 0.5% reported by Zhang et al. (2022c) with Keck / NIRSPEC.  \nZhang et al. (2022c) detected an excess absorption of Ly α in one visit of HD 63433 c atmosphere with HST / STIS. However, we do not detect the evaporating atmosphere of HD 63433 c via H α in the VIS, or the He i triplet in the NIR. Similar situation applies to, e.g. GJ 436 b where Ly α was detected (Kulow et al. 2014), but not H α or He i triplet (Cauley et al. 2017; Nortmann et al. 2018; Villarreal D'Angelo et al. 2021).", '3.8. HD 73583 b': 'We analysed the new TESS data to update the ephemerides on HD73583 planetary system (see App.E) and we used them to analyse the transit of HD 73583 b observed with HARPS-N. The MO spectra was computed without the last three spectra, which show di ff erent values in the serval H α indicator and YABI S-index. The time evolution for the other activity indicators is mainly flat.  \nThe H α residual map and TS are shown in Fig. A.7, where the last spectra clearly exhibit higher flux in the H α line core. The TS shows a small absorption feature at the H α position, which seems to come from the second half of the transit according to the residual map. The signal is -0.46 ± 0.16 % deep (significance < 3 σ ), and has a net shift of -6.1 + 3 . 2 -2 . 5 kms -1 . The priors and posteriors from the fit are shown in Table E.3, and the posterior distributions are plotted in Fig. E.3. However, it is hard to link that absorption to the planet atmosphere for two reasons: i ) the absorption comes mostly from the second half of transit (see the residual map on Fig. A.7), and ii ) the shift of the feature is in the opposite direction of the He i detection from Zhang et al. (2022b). We considered the ∼ 0.5% absorption feature as the upper limit for H α absorption.  \nBecause GIANO-B was not on the telescope the night of the observations, we could not explore the He i triplet. Zhang et al. (2022b) detected an excess absorption of 0.68 ± 0.08 % from one full and one partial transits of HD 73583 b with Keck / NIRSPEC. These observations were re-analysed in Zhang et al. (2023b) reporting a new excess absorption measurement of 0.72 ± 0.08 %. The He i triplet signal has a blue shift, rather than the usual redshift reported in the other He i detections.', '3.9. HD 73583 c': 'As we mention in the previous section, the TESS photometric analysis for the HD 73583 planetary system is shown in App. E. Using the values from Barragán et al. (2022), the uncertainty on HD73583c T c would be 50 minutes, while we now obtained a precision of ∼ 1 min.  \nThe time evolution of the serval H α and dLW indicators show a peak at the beginning of the transit that extends until midtransit. Thus, to compute the TS, we only considered the spectra from the second half of the transit. The other activity indicators show a scattered but flat time evolution. The H α residual map and TS are shown in Fig. A.8, where the H α variability is visible in the residual map. The TS has a small feature at the expected position. We fitted an absorption of -0.54 + 0 . 13 -0 . 14 %, and we present the prior and posterior values in Table E.3, and the posterior dis-  \nbutions are shown in Fig. E.3. Although the fitted absorption is significant ( ∼ 3.8 σ ), we consider that value as an upper limit or a tentative detection. A second transit with better S / N and no stellar activity will confirm these results.  \nThe He i triplet residual map and TS is shown in Fig. A.8. There is no clear evidence of a planetary trace on the residual map or planetary absorption detected on the TS. We placed a 3 σ upper limit to the He i excess absorption of 0.5 %.', '3.10. HD 235088 b': 'We observed one transit of this sub-Neptune-sized planet during this project and its results were analysed in depth in OrellMiquel et al. (2023). We detected a He i blueshifted absorption signal of -0.91 ± 0.11 %, confirming the previous detection from Zhang et al. (2023b) of -0.64 ± 0.06 % absorption. The residual maps and transmission spectra for the H α and He i triplet are shown in Fig. A.9 (adapted from Orell-Miquel et al. 2023).  \nIn this work we adopted the results presented in Orell-Miquel et al. (2023), where an age of 600-800 Myr was estimated for HD235088.', '3.11. K2-77 b': 'Because K2-77 b ephemerides are based on K2 photometry, we decided to analyse the available TESS data and improve the uncertainties on T c for this target (see App. F).  \nAccording to our results, the last spectrum from the CARMENES observations was taken during the egress, missing the post-transit coverage to compute the MO spectrum. Moreover, the pre-transit observations have low S / Ndue to higher airmass.  \nThe di ffi culties to calculate the MO spectrum are visible in the residual maps and TS for both H α and He i triplet lines (Fig. A.10). We could only derive broad 3 σ upper limits for H α and He i excess absorptions of 2.5 % and 2.7 %, respectively.', '3.12. TOI-2046 b': "We set GIARPS mode observations with similar exposure times than for HARPS-N and GIANO-B. Due to the faintness of the host star (J = 10.4 mag), we used GIANO-B's maximum time exposure of 600 s for both spectrographs. The spectra have relatively low S / N around the H α and the He i triplet. Also, the GOFIO pipeline could not properly remove the biggest telluric OH emission peak, and we decided to mask those telluric affected regions. The masked regions are far enough from the He i line to ensure that they do not a ff ect the final results.  \nH α and He i residual maps and TS are shown in Figures A.11. Weonly could derive very broad 3 σ upper limits for H α and He i excess absorptions of 5 % and 2.9 %, respectively.", '3.13. TOI-1807 b': 'We analysed the newly available TESS data to derive updated ephemerids, and refine the T c. The details of the photometric fit are explained in App. G, and our results for the planet properties are shown in Table G.1. However, we planned the observations with older ephemerides (Nardiello et al. 2022), missing part of the transit on 16 December 2022. We observed a second full transit on 23 December 2022, also with better S / N. The H α and He i residual maps and TS for both events are shown in Figures A.12. The second visit does not show evidences of absorption features.  \nWe only used the data from the second transit to compute upper limits due to the di ff erent quality between both observations. We derived a 3 σ upper limit of 0.95 % for H α excess absorption, and an upper limit of 0.80 % for the He i triplet.  \nA transit of this ultra-short period (USP) Earth-like density planet was already observed with IRD, reporting an upper limit to the He i triplet EW of 4 mÅ (Gaidos et al. 2023, ∼ 0.4 % absorption). Because our upper limit is less constraining than the previous observations, we used the results from Gaidos et al. (2023) in the discussion in Section 4.', '3.14. TOI-1136 d': 'One HARPS-N transit of TOI-1136 d was observed to check for the planetary RM signature in the RV time series (Dai et al. 2023a). In this work, we inspected the dataset looking for H α planetary absorption signal. Furthermore, we scheduled a second transit with CARMENES to inspect the whole spectral range. We used nonlinear ephemerides to determine the T c for the second transit (F. Dai, priv. comm.). The RM and CLV e ff ects over the final TS of H α and He i triplet lines are negligible (Fig. H.1). Thus, we did not correct the data from these marginal e ff ects.  \nThe residual maps and TS of the H α for the individual nights are shown in Figure A.13, while the final results are shown in Figure 2. HARPS-N results show a noisy TS with a small absorption feature. However, a similar feature is also detected in the CARMENES H α TS, which has higher S / N. The results from both nights confirm the H α detection on the atmosphere of TOI-1136 d. The combined H α residual map and TS are shown in Fig. A.13. We fitted the signal obtaining an absorption of -1.12 + 0 . 12 -0 . 13 %. Table H.1 presents the priors and posteriors, and Fig. H.2 shows the posterior distributions. Although the signal is a bit deeper in the first half of the transit, the H α transit light curve (Fig. 3) shows an absorption consistent with the transit duration without evidences of tail-like structure.  \nFor the He i triplet, because the CARMENES observations were performed close to the end of the night, the last spectra are a ff ected by solar He i emission. Although we corrected the emission line, some residuals still persisted. Then, we decided to mask the a ff ected spectral range only on particular spectra. The He i residual map and TS are shown in Fig. A.13, and do not show clear evidences of He i absorption. We placed a 3 σ upper limit to the He i triplet excess absorption of ∼ 0.5 %.', '3.15. TOI-1268 b': 'The transit of TOI-1268 b on the 24 of February 2023 was only followed spectroscopically by GIANO-B, and we missed the information about H α , and the stellar variability or activity from the visible wavelength range covered by HARPS-N. MuSCAT2 covered the same transit photometrically with the exposure times initially set to g = 5 s, r = 10 s, i = 10 s, and zs = 15 s and later modified to g = 5 s, r = 10 s, i = 8 s, and zs = 12 s to avoid the saturation of the target star. Our MuSCAT2 transit fit ha a central time of T c = 2460000.66841 ± 0.00014 BJD.  \nThe four bands of the MuSCAT2 photometry do not show the presence of strong stellar activity or the planet crossing in front of any big starspot(s). Moreover, the Paβ , Paγ , and Paδ lines are mainly flat with no correlation with the stellar He i line. Because our observations ended close to the twilight, our last spectra are a ff ected by He i sunlight contamination. Because this kind of emission increases quickly between exposures as the observations are approaching the twilight, it is not well corrected  \nFig. 2: Same as Fig. 1, but for TOI-1136 d. H α results (top panels) are the combination of HARPS-N and CARMENES VIS observations. He i triplet results (bottom panels) are from CARMENES NIR observations.  \n<!-- image -->  \nFig. 3: Transit light curve of H α line from the combined observations of TOI-1136 d. The H α light curve was constructed integrating the counts of the residual map in the planet rest frame around λ 0 using σ (blue) and FWHM (orange) wavelength band passes from Table H.1. The vertical lines represent the di ff erent contacts during the transit.  \n<!-- image -->  \nby the ABAB procedure. We decided to mask those regions affected by the He i emission only in the selected spectra.  \nThe residual map and TS are shown in Fig. 4, where the residual map shows an absorption region at the expected position of the planetary trace. The signal is well detected in the residual map and the TS, where we fitted a blue-shifted absorption of -2.00 + 0 . 15 -0 . 16 % (EW = 19.1 ± 1.9 mÅ). Table I.1 presents the priors and posteriors, and Fig. I.2 shows the posterior distributions. These high-resolution spectroscopy results confirm the He i detection using narrow-band photometry from PérezGonzález et al. (2024). We also constructed the He i triplet transit light curve, shown in Fig. 5. The absorption seems to extend further than the end of the white light transit hinting an He tail, that needs to be confirmed in further observations.', '3.16. TOI-2076 b': 'In this section we present the analyses and results from one CARMENES transit of TOI-2076 b observed in 2023. Previously in May 2022, we missed a transit of TOI-2076 c due to imprecise ephemerides from Hedges et al. (2021). Nowadays, the periods of TOI-2076 planets are well-constrained and refined (Osborn et al. 2022; Zhang et al. 2023b), and we could confirm that none of the three planets were transiting the night of 11 May 2022. However, we take advantage of those observations to study the behaviour of the star between epochs. The values derived from the serval activity indicators for the two datasets are consistent, and there are no di ff erences between epochs. The levels of stellar activity of both visits are very similar and TOI-2076 b transit seems una ff ected by stellar activity o variability. Due to bad weather conditions at the beginning of the night of 13 April 2023, we missed the pre-transit baseline and a very small part of the transit. Moreover, the median S / N of the in-transit spectra is lower ( ∼ 60) than after transit ( ∼ 90).  \nThe TTVs of the TOI-2076 system (Osborn et al. 2022) make the scheduling of its transits complicated and increase the probability of missing part of them. Thus, we complemented our observations on 2023 with a photometric follow-up. We could not perform simultaneous observations on 13 April 2023, but we detected TOI-2076 b ingress one month after (14 May 2023). The details of the photometric analysis and their results are explained in App. J.  \nThe H α residual map and TS are shown in Fig. A.14. The TS close to the H α has consistently negative values that could be considered as a planetary absorption. Although the lack of pre-transit baseline could be the origin of the signal, as for HD63433b transit on 1 November 2021 (Sect. 3.6). Then, a full transit coverage is needed to confirm this tentative signal. We set a 3 σ upper limit of 0.7% to H α exoplanet absorption.  \nThe He i triplet residual map and TS are shown in Fig. A.14. We do not detect clear evidences of He i planetary absorption. We looked for correlations of the He i NIR line with the He i D3  \nFig. 4: Same as Fig. 1, but for TOI-1268 b He i triplet observations with GIANO-B.  \n<!-- image -->  \nFig. 5: Transit light curve of He i triplet from TOI-1268 b detection. The He i light curve was constructed integrating the counts of the residual map in the planet rest frame around λ 0 using σ (blue) and FWHM (orange) wavelength band passes from Table I.1. The vertical lines represent the di ff erent contacts during the transit.  \n<!-- image -->  \n(5877.2 Å), and H Paschen lines, with no conclusive results. We could only put a 3 σ upper limit of ∼ 1% for any He i triplet planetary absorption.  \nTOI-2076 b atmosphere was already targeted looking for He i . The same partial transit was observed from the same mountain with two di ff erent 8-m telescopes with their high-resolution spectrographs. In both cases, a consistent feature was detected but there are di ff erences in the interpretation of the signal. Zhang et al. (2023b) with Keck / NIRSPEC reported a He i detection with an excess absorption of 1.01 ± 0.05 % (EW = 10.0 ± 0.7). On the other hand, Gaidos et al. (2023) fitted a significant excess EWof 8.5 ± 1.4 mÅ, but their analyses did not rule out the stellar origin of the signal. Although CARMENES has the capability to detect such He i excess absorptions (e.g. Salz et al. 2018; AlonsoFloriano et al. 2019; Orell-Miquel et al. 2023), the S / N of our transit only allows us to put an upper limit at ∼ 1%, which is at the same level of the previous detected feature. Further observations in good weather conditions are needed to firmly confirm the TOI-2076 b He i signal. For the purpose of this work, we use the positive results from Zhang et al. (2023b) in the discussion.  \nWe also used the dataset from May 2022 to check for H / He structures around TOI-2076, recently reported in other exoplanets (Zhang et al. 2023c; Gully-Santiago et al. 2024). However, we obtained very similar results when we included those spectra in the analyses of TOI-2076 b H α and He i lines.', '3.17. TOI-1683 b': 'We scheduled three transits of TOI-1683 b for the MOPYS project, one with CARMENES and two with GIARPS. CARMENES transit was stopped due to bad weather and the few taken spectra are useless for atmospheric analyses. The two other transits were observed only with GIANO-B due to problems with the GIARPS mode, and we could not observe until ingress for the first GIANO-B visit (19 September 2023).  \nZhang et al. (2023b) reported an age of 500 ± 150 Myr, based on gyrochronology relations. Here, we derived a new age estimation of 2 + 1 . 3 -0 . 9 Gyr based on gyrochronology, stellar kinematics, and stellar Li i analyses, indicating that TOI-1683 is not a young star. Hence, we do not consider TOI-1683 b as a young planet. The details of our stellar age analyses are explained in Appendix K. Furthermore, we derived new ephemerides for TOI1683 b using the available TESS data from Sectors 19, 43, and 44 in the MAST archive. The photometric analyses and results are shown in Appendix K as well.  \nZhang et al. (2023b) observed one transit in bad weather conditions with Keck / NIRSPEC, and reported a He i detection of 0.84 ± 0.17 %. Although we also observe in poor weather and seeing conditions, we do not detect a significant He i excess absorption in the individual or combined nights from GIANO-B. The residual maps and TS from the individual and combined nights are shown in Fig. A.15. We placed a 3 σ upper limit from our two visits with GIANO-B to the He i excess absorption of 0.7 %, which is ∼ 1 σ consistent with Zhang et al. (2023b) detection. We use the positive detection from Zhang et al. (2023b) in the discussion.', '3.18. WASP-189b': "Stangret et al. (2022) already analysed one HARPS-N and two HARPS observations finding a H α absorption feature of -0.13 ± 0.02 %. However, because the line was a ff ected by the RM e ff ect, they were cautious claiming a detection. Here, we inspected the GIANO-B data taken simultaneously to those HARPS-N observations.  \nThe Si i and He i stellar lines are not detectable in the spectra and the only visible spectral features are the H2O telluric absorption lines, which we masked in the residual map and TS, shown in Fig. A.16. Although the TS has some structure, there is no evidence of He i triplet absorption. We set a 3 σ upper limit of 0.3 % to the He i absorption on WASP-189 b.  \nHAT-P-57b is an 'adolescent' planet (1.00 + 0 . 67 -0 . 51 Gyr; Hartman et al. 2015), where Stangret et al. (2022) found an H α absorption feature ( -0.7 ± 0.2 %). However, the signal is narrow compared to other H α detections from the literature and it is surrounded by pulsations from the host star.  \nWe analysed one archival CARMENES NIR transit to inspect the He i triplet. Although we masked some regions of strong tellurics, the residual map and TS in Fig. A.17 do not show any significant planetary absorption signal. We set a 3 σ upper limit of 1 % to the planetary He i absorption.", '3.20. TOI-2018 b': "The planet candidate TOI-2018 b was initially alerted to the community as a 'young planet', but final analyses of its age were ambiguous and it was not possible to confirm an age below 1 Gyr (2.4 + 0 . 8 -0 . 2 ; Dai et al. 2023b). Although we can not include TOI-2018 b in our sample of young planets, the results obtained from two partial transits observed with GIARPS deserve to be included in this work.  \nThe GIARPS transit on the night of 15 June 2023 was photometrically followed with MuSCAT2, observing the transit in four bands with the exposure times set to g = 15 s, r = 15 s, i = 10 s, and zs = 10 s. Our MuSCAT2 transit fit found a central time of T c = 2459746.4287 ± 0.0021 BJD, confirming that we missed partially the transits in both visits.  \nThe H α residual map and TS from the individual nights and their combination are shown in Figures A.18 and 6, respectively. The TS is flat, and the 3 σ upper limit for H α excess absorption is set to 1.5 %.  \nThe He i triplet residual map and TS from the individual nights and their combination are shown in Figures A.19 and 6, respectively. Although the quality of the two nights is di ff erent, the results of both nights are consistent within 2 σ , and we detect a consistent excess absorption of ∼ 1 % at the expected position of the planetary He i signal. When we combine both nights, we confirm the detection of a red-shifted He i absorption of -1.02 + 0 . 19 -0 . 22 % (EW = 7.8 ± 1.5 mÅ). All the nested sampling material for the individual and combined datasets is shown in App. L. The He i triplet transit light curve (Figure 7) shows an absorption signal consistent with the transit duration, with no evidence of an extended He signal. Although the He i signal is consistent and well detected in the two parcial transits, a full transit observation will help to confirm our detection and the study of TOI-2018 b atmospheric evaporation.", '4.1. Young planets from the literature': "In the framework of the MOPYS project, we adopted the 1 Gyr stellar age as the threshold to classify exoplanets as young ( ≲ 1 Gyr) or old ( ≳ 1 Gyr). We used this nomenclature in the discussion in Section 6. The 1 Gyr threshold is mainly based on the core-powered evaporation timescale, which is longer than the strong photo-evaporation initial stage (X-ray driven).  \nYoung exoplanets are an interesting population that have called the attention of di ff erent research groups. To put the young planet results obtained in this work into context, we inspected  \nthe ExoAtmospheres 5 database looking for other young exoatmospheric analyses. Table 3 presents a compilation of literature high-resolution spectroscopy studies of young transiting exoplanet atmospheres targeting the H α or the He i triplet. 6 Here, we give a short context to those observations from the literature included in Table 3.  \nKELT-9b: As the hottest exoplanet known to date, this planet attracted attention to probe and investigate its extreme atmosphere. Yan & Henning (2018) detected an evaporating exosphere of H α . However, Nortmann et al. (2018, see supplementary material section therein) could only place an upper limit to the He i triplet absorption.  \nK2-25b: The He i triplet was observed using the IRD spectrograph, although the transit was contaminated by telluric OH emission. Gaidos et al. (2020b) reported a 99% confidence upper limit to the transit-associated EW of 17 mÅ.  \nAU Micb: Due to the host star's youth, the H α observations with ESPRESSO were strongly a ff ected by stellar activity, making it impossible to set an upper limit (Palle et al. 2020b). Recently, Rockcli ff e et al. (2023) reported a Ly α detection but only in one of the two visits with HST / STIS. IRD and Keck / NIRSPEC spectrographs obtained a 99% confidence upper limit to the He i EW of 4.4 and 3.7 mÅ, respectively (Hirano et al. 2020). Allart et al. (2023) observed one transit with SPIRou spectrograph finding a significant He i absorption feature (0.37 ± 0.09 %). However, the authors finally reported a conservative 3 σ upper limit of 0.26 %, consistent with previous observations.  \nK2-136c: Gaidos et al. (2021) put a 99% confidence upper limit to the He i triplet EW of 25 mÅ with one transit observed with the IRD spectrograph. The He i lines position in between the telluric OH emission and H2O absorption lines complicated the calculation of the transmission spectra.  \nDS Tucb: The H α was analysed with ESPRESSO and HARPS spectrographs observations. However, the transits were a ff ected by stellar activity and an H α upper limit could not be set (Benatti et al. 2021).  \nV1298 Tau system: This 20-Myr old multi-planet system has seen di ff erent attempts to study the presence of He i in their planetary atmospheres. Gaidos et al. (2022) used the IRD spectrograph to observe, during di ff erent nights, the star alone and a transit of planet b. An increasing He i absorption was detected during the transit, but the authors proposed other explanations besides the planetary absorption. Using the narrowband helium filter technique, Vissapragada et al. (2021) observed the transits of planets b and d. For planet b they do not require extra absorption to explain the flux decrease, while they found a tentative excess absorption combining two partial transits of planet d, but it requires a significant transit time variation.  \nWASP-52b: The planet was observed with ESPRESSO detecting H α , and other atomic species as well (Chen et al. 2020). The presence of He i in its atmosphere has been studied, first with a tentative detection using narrow-band photometry (Vissapragada et al. 2020, 2022a) and later confirmed with the Keck / NIRSPEC spectrograph (Kirk et al. 2022). However, a recent paper by Allart et al. (2023) did not find He i absorption in two visits with the SPIRou spectrograph, and derived a 3 σ upper limit of 1.69 %.  \nFig. 6: Same as Fig. 1, but for TOI-2018 b combined observations of H α with HARPS-N and He i triplet with GIANO-B.  \n<!-- image -->  \nFig. 7: Transit light curve of He i triplet from the combined nights of TOI-2018 b. The He i light curve was constructed integrating the counts of the residual map in the planet rest frame around λ 0 using σ (blue) and FWHM (orange) wavelength band passes from Table L.3. The vertical lines represent the di ff erent contacts during the observation.  \n<!-- image -->  \nHAT-P-70b: The atmosphere of this ultra-hot Jupiter was inspected by Bello-Arufe et al. (2022) using HARPS-N. They detected absorption coming from the H α , H β , and H γ lines, and many other atomic and molecular species as well.  \nWASP-80b: Its atmosphere has been targeted in search of He i absorption but without success (Fossati et al. 2022; Allart et al. 2023). Salz et al. (2015) is the only reference which provides a formerly calculated age of < 200 Myr. However, the observed X-ray luminosity (log L X ∼ 27.8 erg s -1 in Salz et al. 2015, down to 27.5 erg s -1 in Sanz-Forcada et al. in prep.) yields a log L X / L bol ∼-5 . 0, implying an age of ∼ 3 Gyr using SanzForcada et al. (2011) ageL X relations. Moreover, the TESS light curve do not suggest a rotational period of ≲ 15 d. Therefore, WASP-80 is probably older than 1 Gry according to the gyrochronology method shown in Figure K.3 ( G -J = 2.05).", '4.2. Older planets from the literature': 'To put in context the young exoplanet He i (planets with ages ≲ 1 Gyr) findings, we complemented them with old exoplanet high-resolution He i observations (planets with age ≳ 1 Gyr) from the literature. We have compiled a He i triplet database, detailed in Table M.1, which we constructed using the ExoAtmospheres database and literature results. For consistency, we only considered He i triplet results from high-resolution spectrographs. That is, the narrow-band photometry detections of HATP-26 b (Vissapragada et al. 2022b) and TOI-1420 b (Vissapragada et al. 2024a) were not included. 7 Although we analysed the H α and we cited some results on the Ly α as well, Table M.1 only reports the H (H α or Ly α ) observations for the planets with He i triplet observations, except for the young hot Jupiter HATP-70 b with H α detection (Bello-Arufe et al. 2022). The young Neptune DS Tuc b is not included in Table M.1 because no upper limit could be set from H α observations (Benatti et al. 2021) and no He i triplet observations could be found in the literature.  \nGiven the number of planets and details of each observation, we do not discuss the planets individually here. We refer to Table M.1 and the appropriate references.', '4.3. A note on He I variability': 'The planetary He i triplet is known to show variability in its strength but also in its detectability (e.g. Palle et al. 2020a; Zhang et al. 2022a). In this work, we presented two upper limits for TOI-2076 b and TO1683 b, for which Zhang et al. (2023b) reported He i detections. Our upper limits are consistent at 1 σ with Zhang et al. (2023b) detections. Thus, He i triplet variability may be one reason for our absence of planetary signals. Variability was also invoked by Orell-Miquel et al. (2023) to explain the ∼ 2 σ deeper absorption found in HD 235088 b. We note that Gaidos et al. (2023) derived non-conclusive results for TOI-2076 b  \nTable 3: Compilation of H α and He i observations from high-resolution spectroscopy facilities of young ( ≲ 1 Gyr) transiting exoplanets, from the literature and this work (Sect. 3). A cross ( × ) indicates when the stellar activity prevented to measure any planetary absorption or derive upper limits.  \nReferences. (1) Casasayas-Barris et al. (2019). (2) EW at 99% confidence of 5.7 mÅ (Gaidos et al. 2020a). (3) Stangret et al. (2021). (4) Zhang et al. (2022c). (5) Zhang et al. (2022b). (6) Zhang et al. (2023b). (7) Orell-Miquel et al. (2023). (8) EW at 99% confidence of 4 mÅ (Gaidos et al. 2023). (9) Stangret et al. (2022). (10) Yan & Henning (2018). (11) Nortmann et al. (2018). (12) EWat 99% confidence of 17 mÅ (Gaidos et al. 2020b). (13) Palle et al. (2020b) (14) EWat 99% confidence of 3.7 mÅ (Hirano et al. 2020). (15) EWat 99% confidence of 25 mÅ (Gaidos et al. 2021). (16) Benatti et al. (2021). (17) Chen et al. (2020). (18) Kirk et al. (2022). (19) Bello-Arufe et al. (2022). (20) Gaidos et al. (2022). (21) Fossati et al. (2022)  \nNotes. ( † ) TOI-2018 b and TOI-1683 b are included in the table, but they are likely not young planets.  \nHe i signal, and TOI-1683 b transit from Zhang et al. (2023b) was performed in poor observing conditions. We want to stress the importance of re-observing targets as a sanity check to confirm previous results (detections or upper limits), but also to study the planetary He i triplet variability.  \nOn the other hand, Krolikowski et al. (2024) analysed the variability of the He i triplet of young stars. In particular, V1298Tau, K2-100, K2-136, K2-77, HD 63433, and TOI-2048 were in the list of observed stars. They found that young stars show higher variability, with V1298 Tau being extremely variable, which could explain the non-conclusive results from the He i observations. The stellar He i variability decreases rapidly and keeps constant for stars older than ∼ 300 Myr (Krolikowski et al. 2024). Although the stellar He i triplet line is variable, the short-term variability might not have a significant impact on the transit observations performed during the same night (Fuhrmeister et al. 2020; Krolikowski et al. 2024). In fact, H α is more sensitive to stellar variability than the He i triplet (Fuhrmeister et al. 2020) and the analyses of those lines on AU Mic b and DS Tuc b are good examples (Palle et al. 2020b; Hirano et al. 2020; Benatti et al. 2021).', '5. Criteria for detections, non-detections, and non-conclusive measurements': 'A problem we encountered when studying the He i signal (and H α as well) was how to deal with the observations where only an upper limit value is given, which are the majority of the cases (52 of 69 in Table M.1). In this work, we classified the upper limit measurements in two groups: a ) non-conclusive measurements (i.e. the upper limit value is large, and it does not actually constrain the presence of the atom in the exoplanet atmosphere) and b ) non-detection (i.e. the upper limit value is low enough to confidently assume there is no significant presence of that atom in the exoplanet atmosphere).  \nTo assign observations to either category, for the He observations, we used the relationship that the observed and theoretical mass-loss rates ( ˙ m obs and ˙ m theory, respectively) seem to follow, represented in Figure 8. We present ˙ m obs and ˙ m theory equations in Sect. 6.3, and we further explore and discuss the trend in that section. For the time being, the fitted line between ˙ m obs- ˙ m theory seems to account for the di ff erences between planetary characteristics, and populations (Vissapragada et al. 2022b; Zhang et al. 2023a, and Sect. 6.3). Therefore, we consider that relation  \nFigure 10 is similar to Figure 9 but focuses only on the He i triplet, and presents a comparison between the He i EW signals and the stellar age. The planets are marked and colour-coded according to their He i results. Allart et al. (2023) reported that a trend between He i and stellar age was noticeable in their sample of eleven planets, however the authors refrained from further conclusions due to their small sample size and the lack of precise stellar ages. Although in this work we used a larger sample (18 young and 35 old planets, including detections and nondetections, but less homogeneous than the Allart et al. 2023 one), we do not notice a clear correlation between He i EWand stellar age.  \n<!-- image -->  \nto classify the He i upper limits. Although other criteria might be chosen, it is used consistently for the whole sample.  \nWe computed the ˙ m obs upper limits from the He i EW upper limit values. If the ˙ m obs upper limit falls above the fitted line, the ˙ m obs upper limit value is higher than the a priori expected signal, and so it is not constrained. Thus, we refer to those measurements as non-conclusive (orange down-pointing triangles in Fig. 8). On the other side, the ˙ m obs upper limits falling below the line can be considered as non-detections (red crosses in Fig. 8). We also considered as non-detections the ˙ m obs upper limits that fall within 1 σ of the fit (grey shaded are in Fig. 8).  \nWe note that F XUV is unknown for some of the exoplanets in our sample. For those planets with no measured F XUV, we performed a similar procedure but in a 1 /ρ XUV versus ˙ m obs diagram (not shown), which is a proxy for ˙ m theory versus ˙ m obs . ρ XUV is the planet density when considering as the radius its XUV radius, which is di ff erent from the planet density ρ p (see Sect. 6.3).  \nWe did some exceptions when classifying some He i observations. GJ 1214 b: we re-classified as a non-detection the tentative He i signal reported by Orell-Miquel et al. (2022) due to other upper limits reported in the literature (Petit dit de la Roche et al. 2020; Kasper et al. 2020; Spake et al. 2022; Allart et al. 2023). V1298 Tau b: because of the unclear origin of the signal detected by Gaidos et al. (2022), we set this observation as nonconclusive. V1298 Tau c is in the same situation as planet b, and its He i upper limit is non-conclusive. Moreover, the measurements for the V1298 Tau planets are consistent with the stellar He i line variability range derived for their host star (Krolikowski et al. 2024). WASP-76 b: its He i feature was presented as upper limit in Casasayas-Barris et al. (2021a) but here we consider it as a non-detection due to the planet position well below the ˙ m obs-˙ m theory line.  \nFor H α , we simply considered the upper limits as nondetections. However, if the H α upper limit absorption is comparable to that of He i , we used the He i classification for both measurements. We only considered few Ly α detections for some particular exoplanets. The classification and nomenclature described in this section (detection, non-detection, and non-conclusive) is used all across the manuscript.  \nWe note that absoprtion measurements of the three lines (Ly α , H α , and He i triplet) would be the ideal case to determine whether an exoplanet is undergoing strong atmospheric escape. However, these observations are not available for all targets,  \nFig. 8. Relationship between the observed ( ˙ m obs) and the theoretical ( ˙ m theory) energylimited mass-loss rates. We define XUV until the He i ionisation range ( λ = 5 - 504 Å). The black line indicates the fitted relationship (shown in the legend) and the shaded area the 1 σ uncertainty. He i observations are coded as blue circles for detections, red crosses for nondetections, orange down-pointing triangles for non-conclusive. We did not plot the error bars of ˙ m theory due to the very large uncertainties associated with the F XUV values and its calculation. Every planet has its name labelled: in black for detections and the others in grey. The two new detections presented in this work, TOI2018 b and TOI-1268 b, are in good agreement with the predicted trend.  \nand may be hard (even impossible) to obtain for some individual stars (e.g. Ly α interstellar medium extinction, H α variability in active stars, or low He i triplet population). When possible, we consider the results from the three lines to determine atmospheric escape detections.', '6. Discussion': 'We have organised the discussion of our results in the following manner: Section 6.1 gives a general view of the evaporation tracers across stellar age, planet radius, period, and mass. In Sect. 6.2 we constrain the cosmic shoreline from He i detections. In Section 6.3 we explore the relation between He i detections and the energy-limited mass-loss rates. Finally, we explore further relations between He i detections and di ff erent planet and stellar properties in Sections 6.4, 6.5, and 6.6.', '6.1. Evaporation tracers of planetary atmospheres across stellar age': 'Figure 9 displays the planetary radius versus stellar age, with the planets colour-coded according to their evaporation measurements. We considered the He i triplet and / or H α absorption detections (and Ly α in some particular cases) as proxies of evaporation signs. Thus, an evaporation detection means that either He i or H α or both have been positively detected.  \nThe first 100 Myr are critical for the planetary atmospheric evolution, according to photo-evaporation models (see Introduction in Sect. 1). Unfortunately, there are only a few planets in Figures 9 and 10 with ages ≲ 100 Myr. The youngest objects have ages of ∼ 20 Myr (AU Mic b and V1298 Tau c & b) and  \nFig. 9: Planetary radius vs stellar age diagram of evaporation (H and He i ) detections (blue circles), non-detections (red crosses), and non-conclusive observations (orange down-pointing triangles). Planetary radii and stellar ages are from Table M.1. The top panel shows the summed histogram of evaporation detections (blue line) and non-detections (red line) across stellar age. The grey points represent all known planets with radius and age determined with precision better than 30% and 50%, respectively (data from NASA Exoplanet Archive).  \n<!-- image -->  \nFig. 10: Equivalent width (EW) of He i detections (blue circles), non-detections (red crosses), and non-conclusive observations (orange down-pointing triangles) as functions of stellar age. The top panel shows the summed histogram of He i detections (blue line) and non-detections (red line) across stellar age bins.  \n<!-- image -->  \nthen there is a lack of He i and H α observations until ∼ 100 Myr (WASP-80b, and K2-77 b). WASP-80 b, with a formerly calculated age of < 200 Myr (Salz et al. 2015), has no detection of He i triplet. However, a possible older age of WASP-80 (see Section 4.1) and the calculated log L XUV at the planet separation (Table M.1) could explain the lack of detection of the He i triplet. AUMicb non-detection is harder to reconcile, but di ff erent scenarios can be invoked (e.g. stellar activity masking possible detections or H / He ratio). He i and H α observations of V1298 Tau b &c resulted in non-conclusive measurements, mainly due to the strong stellar activity levels of the young host star. However, the masses derived for planets b and e indicate that their densities are similar to older planets, suggesting they are not inflated and that they contracted faster than expected (Suárez Mascareño et al. 2021).  \nWhile the non-detections of WASP-80b, AUMicb, V1298Taub & c do not seem to fit the predictions from  \nphoto-evaporation models (Lopez et al. 2012; Owen & Jackson 2012; Owen & Wu 2013, 2017; Owen & Lai 2018; Dawson & Johnson 2018), the sample of ≲ 100 Myr-old planets is too small to draw conclusions. A larger sample of ≲ 150-Myr-old planets is needed to explore the photo-evaporation timescales.  \nFrom Fig. 9 and 10, the youngest planet with He i detection is TOI-1268 b (245 ± 135 Myr), and the youngest planet with H α detection is MASCARA-2 b (200 + 100 -50 Myr). The number of planets with evaporation detections increases until peaking at ∼ 2 Gyr, and then it is roughly constant, but this is related to the age distribution in our sample. The proportion between detections and non-detections remains constant (within error bars) with age, except at old ages ( > 5 Gyr) when non-detections dominate (see Figs. 9, and 10). This non-detections domination is more pronounced when considering only the He i observations (Fig. 10), and extended over all stellar ages, except over the 1-3 Gyr range.  \nOur results on evaporation timescales agree with the conclusions from Loyd et al. (2020) on the radius gap, and Christiansen et al. (2023) on the hot Neptune population (see Fig. 6 in their work), and photo-evaporation mechanism is not more supported than the core-powered one. The fact that we do not see a decrease in the evaporation tracers, until very old ages ( ∼ 5 Gyr, see Figs. 9, and 10), might be marginally more consistent with the core-powered timescale of Gyr (Ginzburg et al. 2016, 2018; Gupta & Schlichting 2020).', '6.1.1. Radius-period diagram across stellar age': 'Figure 11 compares the population of young ( ≤ 1 Gyr, top panels) and old ( > 1 Gyr, bottom panels) planets in a radius-period diagrams for only He i triplet observations, only H (mainly H α with few Ly α ) observations, and both evaporation proxies. Figures 9, 10, and 11 show no clear di ff erences for the evaporation of the gas giant planets before and after 1 Gyr. The detections of evaporation are evenly spread over the stellar ages, with more preference of H α detection than from He i for the young gas giants. We note that there are no evaporation detections of young and old planets below the radius gap or for Earth-like planets, supporting rocky planets are not under extreme atmospheric massloss processes, at least after ∼ 300 Myr which is the age of the youngest rocky planet in our dataset (TOI-1807 b).  \nAllan et al. (2024) simulated how the He i triplet planetary signal from a highly irradiated ( a p = 0.045 AU) gas giant planet ( M p = 0.3 MJ) around a K-dwarf star changes with the stellar age. They assumed a typical H / He ratio of 98 / 2 (Lampón et al. 2020, 2021b; Orell-Miquel et al. 2023). From their simulations, they derive excess absorption peaks of 4-7 % for young (16-550 Myr) Hot Jupiters, while at 5 Gyr the excess absorption would be of ∼ 1.5%. They state that a close-in ( a p < 0.1 AU) planet with a radius of 1-2 RJ transiting a < 150-Myr-old K dwarf star would be the best target to test their evolution models. Although the He i triplet detection is favoured by the extreme radiation in XUV range ( λ < 504 nm; Sanz-Forcada & Dupree 2008) and the TS simulations from Allan et al. (2024), none of the three 20-Myr-old planets analysed here present a clear detection of He i . WASP-80 b ( a p = 0.035 AU, R p = 1 RJ, M p = 0.5 MJ, < 200 Myr, spectral type ∼ K7V) is the closest planet from Table M.1 to the simulated one in Allan et al. (2024) but it has very low He i upper limits (Fossati et al. 2022; Allart et al. 2023). WASP-52 b might be in agreement with the simulations, although the largest He i excess absorptions to date come from older planets (e.g. HAT-P-67 b, Gully-Santiago et al. 2024; HATP-32 b, Czesla et al. 2022; and WASP-107 b, Kirk et al. 2020). In Fig. 11 WASP-52 b is the only young gas giant planet with a He i  \nFig. 11: Radius-period diagrams for planets with ages ≤ 1 Gyr (top panels) and > 1 Gyr (bottom panels) with He i (left panels) and H (H α or Ly α , middle panels) observations, and evaporation (combination of He i and / or H observations, right panels). Detections and non-detections are marked as blue circles and red crosses, respectively. The radius gap (Van Eylen et al. 2018) is marked as a dashed green line. The grey points represent all known planets with period and radius determined with a precision better than 25% (data from NASA Exoplanet Archive).  \n<!-- image -->  \ndetection. The ratio of He i detections to non-detections seems larger for old rather than young planets, but there are fewer differences when comparing the H detections to non-detections.', '6.1.2. Small planet evaporation across stellar age': 'Figure 12 shows the radius-mass diagram for small planets ( R p < 5 R ⊕ and M p < 30 M ⊕ ), also comparing evaporation proxies for young and old planets. Planets with R p ∼ 1.5-3 R ⊕ fall in a degenerated region of the mass-radius diagram, and their bulk compositions can be consistent with a large range of models, from water worlds (planets with a large water mass fraction) to planets with rocky cores with H / He envelopes (Zeng et al. 2019). For these planets, in Figure 12, we found a mixture of detections and non-detections, with no di ff erence between young and old planets.  \nPlanets with R p > 3 R ⊕ are well above the water-rich composition line and are supposed to be gaseous with very light envelopes and low densities (Luque & Pallé 2022). In our sample there are four young and five old pu ff y planets with evaporation observations (see right panels of Fig. 12). While only 1 in 4 young planets has an evaporation detection, 3 out of 5 old planets have a detection, hinting that atmospheric escape of pu ff y subNeptunes is stronger at ages older than ∼ 1 Gyr. Still, the numbers are small, and a larger sample is needed to confirm these findings.  \nFor the planets above the radius valley, Malsky et al. (2023) predicted a He enhancement due to the di ff usive separation of the atmospheric constituents where the He and metals are preferentially retained while H is evaporated. The timescale of this mechanism is comparable to the planet lifetime ( ∼ 10 Gyr) and planets with ≲ 1 Gyr have not had time to show the e ff ects of favoured H evaporation (Malsky et al. 2023). So, our young planet population is too young to su ff er this di ff erential escape process, and even some of our old planets could be considered young as well. Then, the long timescale might explain why Figure 12 does not show more He i detections for the old waterworlds than for the young ones. Malsky et al. (2023) predict that TOI-1235 b will have an He enhanced atmosphere at 10 Gyr, but Krishnamurthy et al. (2023) put a very restricted He i upper limit, although its age is poorly constrained (5 + 5 -4 . 4 Gyr). The other planets listed in Malsky et al. (2023, Table 1) have no H / He observations to test their predictions.  \nAlthough Malsky et al. (2023) focus only on the small planet population ( ≲ 3 R ⊕ ), the di ff usive separation and the subsequent preferential H evaporation could be the mechanism to explain some observational results from Fig. 11 and 12 regarding large and intermediate-sized planets. i) There are no He i detections in young pu ff y planets but there are for old ones. TOI-1136 d is the only young pu ff y planet with H α detection, and the only three old pu ff y planets with H observations both show H evaporation. ii) WASP-52 b is the only young gas giant with He i detection (and also has H α detection) while there are several detections on  \nFig. 12: Radius-mass diagrams for planets with ages ≤ 1 Gyr (top panels) and > 1 Gyr (bottom panels) with He i (left panels) and H (H α or Ly α , middle panels) observations, and evaporation (combination of He i and / or H observations, right panels). Detections and non-detections are marked as blue circles and red crosses, respectively. The solid lines are the theoretical models from Zeng et al. (2019) for Earth-like (green) and 50% H2O + 50%rocky (blue) compositions. The grey points represent all known planets with mass and radius determined with a precision better than 20% (data from NASA Exoplanet Archive).  \n<!-- image -->  \nold gas giants, and H α is extensively detected in young and old gas giant planets.', '6.2. Cosmic shoreline from He i observations': "The cosmic shoreline (Zahnle 1998) is an empirical division found in the Solar System bodies that splits them between those that retain a certain amount of atmosphere and those that are purely bare rocks without atmosphere. The division was extended to the extrasolar planet population with success by Zahnle & Catling (2017). Determining precisely the cosmic shoreline is important as it has often been used to predict the existence of an atmosphere for newly discovered planets, and to argue for atmospheric characterisation follow-up (e.g. with the James Webb Space Telescope). Observationally, extended atmospheres are relatively easy to detect via H and He i absorptions. Here, we use these observations to better constrain the cosmic shoreline.  \nZahnle & Catling (2017) found a power law between planet bolometric instellation ( I p) and planet velocity escape v esc that follows I p ∝ v esc 4 . However, the radiation that shapes the exoplanet atmosphere is the X-rays and EUV stellar flux received by the planet during its early stages. Zahnle & Catling (2017) assumed the approximation of X-ray luminosity saturation (Jackson et al. 2012) to estimate the total extreme radiation (X + EUV radiation) received by the planet. Following Zahnle & Catling (2017, Eq. 27), we compute the cumulative X-rays and EUV in-  \nellation ( I X + EUV; considering λ < 100 nm) as a function of the planet instellation ( I p) and the stellar bolometric luminosity ( L ⋆ ):  \nI X + EUV = I p ( L ⋆/ L ⊙ ) -0 . 6 . (2)  \nWe computed v esc as  \nv esc = q 2 GM p / R p , (3)  \nwhere G is the universal gravitational constant. The empirical relation found between I X + EUV with v esc also follows the same power law ( I X + EUV ∝ v esc 4 ), as for I p.  \nZahnle & Catling (2017) presented the relationships between I p and v esc, and between I X + EUV and v esc for the planets and small bodies of the Solar System along with the exoplanet population. To test the predictions of the cosmic shoreline, we compared the proposed empirical relationships to the actual He i detections. Figure 13 reproduces Figures 1 and 2 from Zahnle & Catling (2017), and includes also the He i observations. Because the cosmic shorelines are empirical relations, only the slope in logarithmic scale is determined. Thus, we get from Zahnle & Catling (2017, Figs. 1 and 2) the independent terms to plot the empirical equations.  \nBy definition, all He i (and H α ) detections must be below the cosmic shorelines plotted in Fig. 13. However, we find seven gas planets with He i (or H α ) detections in the region of supposedly  \nFig. 13: Cosmic shoreline compared to He i observations. He i non-detections are denoted by red error bars with no marker. He i detections are circles scaled by their Rp and colour-coded by their stellar age (lateral colour bar). Instellation relative to Earth vs escape velocity ( left panel ) and cumulative instellation in the XUV range relative to Earth vs escape velocity ( right panel ) graphics. We marked the cosmic shoreline, as I ∝ v esc 4 , from Zahnle & Catling (2017, solid green line) and constrained from the observations (dashed black line). Planets sitting above or on the cosmic shore and with H / He detections are labelled.  \n<!-- image -->  \nbare rocky planets. Those are the ones labelled in Fig. 13 panels. In Fig. 13 left panel, HAT-P-32 b and HAT-P-67 b are clearly located above the line, although their 1 σ uncertainties fall in the limit of the cosmic shoreline. TOI-1136 d, WASP-76 b, and KELT-9b which have He i non-detections but H α detections, are also over or above the cosmic shoreline. Moreover, in Fig. 13 right panel, HAT-P-32 b and WASP-107 b fall over the line, and GJ3470b and TOI-1136 d are clearly above the cosmic shoreline for energy-limited regime. With the exception of TOI-1136 d, the planets that contradict the shoreline are not young planets, with stellar ages > 1 Gyr.  \nZahnle & Catling (2017) stated that the I ∝ v esc 4 and I X + EUV ∝ v esc 4 lines are 'drawn in by hand to guide the eye'. Thus, the evaporation detections above the shoreline can help to constrain the independent terms of those by-hand equations. For that purpose, we considered the planet whose uncertainties have the largest separation from the by-hand line to compute the limits of the cosmic shoreline. We took the coordinates of the extreme uncertainty as a point to calculate the line equation. We took HAT-P-67 b for the I ∝ v esc 4 line, and GJ 3470 b for I X + EUV ∝ v esc 4 although very similar results were obtained with TOI-1136 d. Then, it is trivial to get the line equation with the slope and one point. The cosmic shoreline equations constrained from the evaporation detections are  \nlog( I / I ⊕ ) = 4 log( v esc [km s -1 ]) -2 . 04 (4)  \nand  \nlog( I X + EUV / I X + EUV ⊕ ) = 4 log( v esc [km s -1 ]) -2 . 51 . (5)  \nThey are shown in Figure 13 left and right panels as dashed black lines, respectively. In both cases, the cosmic shoreline moved to higher radiation levels reducing the amount of exoplanets with no atmosphere. Further atmospheric observations and more precise measurements of the planets close to the shoreline will allow a better constraint of the observed cosmic shoreline.  \nFig. 14: He i transmission signal strength for the planets with detections (blue circles with error bars in black) and non-detections (red down-pointing triangles in grey) as a function of the stellar F XUV ( λ = 5 - 504 Å) at the planet distance. We show the equivalent height of the He i atmosphere, δ Rp , normalised by the atmospheric scale height, H eq. Data from Table M.1 and references therein.  \n<!-- image -->", '6.3. Observed versus theoretical mass-loss rates': "Since the first He i detection in planetary atmospheres, there have been many attempts to predict the presence of this line using other observable planetary parameters. Nortmann et al. (2018) performed one of the first attempts by computing δ Rp / H eq. That is: the equivalent height of the He i absorbing atmosphere δ Rp = ( R 2 p + D He R 2 ⋆ ) 1 / 2 -R p, where D He is the He i absorption peak in %, divided by the atmospheric scale height  \nH eq = ( k B T eq) / ( µ g p), where k B, µ , and g p are the Boltzmann constant, the mean molecular mass, and planet gravity, respectively. Figure 4 in Nortmann et al. (2018) presented a correlation between the He i absorption signal and F XUV ( λ = 5 - 504 Å), but with a sample of only five planets. Figure 14 reproduces the same diagram with the current set of 43 He i observations (detections and non-detections with F XUV measurements). The trend is not that clear and there is no evident pattern between detections and non-detections, maybe due to the non-uniformity of the sample (e.g. di ff erent instruments and data analysis techniques).  \nRather than using the strength of the He i absorption, Vissapragada et al. (2022b), and later Zhang et al. (2023a), explored the relation between the mass-loss rates derived from the observations ( ˙ m obs) as a function of F XUV / ρ p ( ρ p is the planet density). While Vissapragada et al. (2022b) derived the observed mass-loss rates comparing the measured excess transit absorptions with a grid of Parker wind models, Zhang et al. (2023a) estimated the observed mass-loss rates from an order-ofmagnitude method. Furthermore, Zhang et al. (2023a) found that the theoretical maximum energy-limited mass-loss rate ( ˙ m theory) is proportional to F XUV / ρ XUV ( ρ XUV is defined below). However, they reported very similar results when using F XUV / ρ p. Despite those di ff erences, both works found a positive correlation between ˙ m obs and F XUV / ρ p, and, their energy-limited outflow e ffi -ciencies agree within uncertainties (Vissapragada et al. (2022b): η 0 = 0.41 + 0 . 16 -0 13 ; Zhang et al. (2023a): η 0 = 0.31 ± 0.06).  \n. Here, we focus on the relationship between ˙ m obs and ˙ m theory while adding the two new He i detections presented in this work (TOI-1268 b and TOI-2018 b), which is the main di ff erence with the study performed by Zhang et al. (2023a). For consistency, we followed the indications from Zhang et al. (2023b,a) to compute the ˙ m obs- ˙ m theory diagram with our set of He i observations.  \nWe made the same assumptions as in Zhang et al. (2023b,a) to calculate ˙ m obs with the order-of-magnitude method as  \n˙ m obs = m e m He c s c 2 EWR ⋆ 0 . 25 f e 2 λ 2 0 Σ gl fl ∝ EW · R ⋆. (6)  \nEquation 6 is computed in cgs units, where m e is the electron mass, m He is the He atomic mass, c s is the sound speed (assumed to be c s = 10 km s -1 ), c is the speed of light, EW is the equivalent width of the He i signal, R ⋆ is the stellar radius, 0.25 comes from assuming that 25 % of mass out-flow is He (atoms or ions), f is the fraction of He atoms in the metastable ground state (assumed to be f = 10 -6 ), e is the electron charge, λ 0 is the He i wavelength (fixed to 10833.3 Å), Σ gl fl is the sum of the product of the degeneracy and oscillator strength over the three lines of the He i triplet which is 1.62 (Zhang et al. 2023a).  \nThe theoretical maximum energy-limited mass-loss rate ( ˙ m theory; Caldiroli et al. 2022; Zhang et al. 2023a), assuming that all the XUV flux received by the planet is spent on evaporation, is computed as  \n˙ m theory = π R 3 XUV F XUV G M p = 3 F XUV 4 G ρ XUV ∝ F XUV /ρ XUV , (7)  \nwhere F XUV is the XUV flux ( λ = 5 - 504 Å), R XUV is the planetary XUV photosphere radius, G is the gravitational constant, M p is the planet mass, and ρ XUV is the planet density using the R XUV. We followed the equations and indications from Wang & Dai (2018); Zhang et al. (2023b,a) to estimate R XUV as  \nR XUV = R p 1 + β -1 ln ( ρ atm . XUV /ρ atm . phot) , (8)  \nwhere  \nρ atm . phot = P µ k B T eq ; β ≡ G M c µ R p k B T eq . (9)  \nIn Eqs. 8 and 9, P is the pressure at the white-light planet radius (assumed P = 100 mbar), µ is the (dimensional) mean molecular mass (we assumed µ = 1.3 × m H; m H is the H atomic mass), k B is Boltzmann constant, T eq is the equilibrium temperature, M c is the planet mass core, and ρ atm . XUV is the planet atmosphere's density at the R XUV layer. We approximated M c ≃ M p and used ρ atm . XUV = 10 -15 g cm -3 . Our computed values for ρ XUV are consistent within errors with those presented in Zhang et al. (2023a, Table 3).  \nFigure 8 presents our computed observed and theoretical mass-loss rates, recovering the predicted positive correlation for the He i detections. For a quantitative analysis, we fitted a power law as ˙ m obs = η 0 ( ˙ m theory) α . We used scipy 's orthogonal distance regression to fit the equation log 10 ( ˙ m obs) = α log 10 ( ˙ m theory) + log 10 ( η 0). We obtained an energy-limited e ffi ciency of η 0 = 0.35 ± 0.04, which is consistent with the previous derived values (Vissapragada et al. 2022b: η 0 = 0.41 + 0 . 16 -0 . 13 ; Zhang et al. 2023a: η 0 = 0.31 ± 0.06). TOI-1268 b and TOI-2018 b positions in the diagram agree with the previous detections, and they are over the fitted line. For the linearity between log mass-loss rates, we got α = 0.43 ± 0.05, also consistent with Zhang et al. (2023a, α = 0.50 ± 0.08). When using ρ p instead of ρ XUV, we get slightly larger uncertainties for η 0 and α but consistent within uncertainties ( η 0 = 0.39 ± 0.06, α = 0.42 ± 0.06).  \nIf the evaporation on those planets was in the energy-limited regime, they should exhibit a linear relation between ˙ m obs and ˙ m theory. Although we computed the mass-loss rates as an orderof-magnitude and with some approximations, α is ∼ 11 σ apart from 1. We find a sub-linearity relation ( α ∼ 0 . 43 < 1) and the e ffi ciency of the ˙ m obs decreases with ˙ m theory, as already pointed out by both Vissapragada et al. (2021) and Zhang et al. (2023a). However, this is in agreement with the theory as models predict that as the F XUV increases, the photo-evaporation escape begins to lose e ffi ciency via radiative cooling (e.g. Murray-Clay et al. 2009; Caldiroli et al. 2022). The H / He ratio or di ff erent hydrodynamical regimes (e.g. Lampón et al. 2020, 2021b,a, 2023) may also contribute to derive a sub-linear relation.  \nAlong the ˙ m obs- ˙ m theory sub-linearity, five planets (TOI2134 b, TOI-2018 b, HD 235088 b, HAT-P-11 b, and HAT-P18 b) show an unphysical > 100 % photo-evaporation e ffi ciency. ˙ m theory, as it is computed in Eq. 7, assumes the photo-evaporation scenario. However, there are other mechanisms that could enhance the atmospheric mass-loss, namely the core-powered and the Roche lobe overflow. The core-powered mechanism may contribute in the atmospheric escape of those planets as they are relatively young (except HAT-P-11 b and HAT-P-18 b). However, a core-powered contribution can not be confidently assumed in those planets, as other parameters as the H / He ratio, the stellar flux or the atmospheric heating e ffi ciency may contribute to change the derived ˙ m obs- ˙ m theory relation and yield unrealistic e ffi ciencies. On the other hand, Roche lobe overflow is not expected, as these planets are not extremely close-in. Nonthermal processes (e.g. ion pick-up or sputtering) may also help to achieving a > 100 % e ffi ciency in those planets (Güdel et al. 2014).  \n<!-- image -->", '6.4. He i dependence with the Hill radius': "Because the F XUV is a critical factor for the detectability of extended He atmospheres via the He i triplet in the NIR (SanzForcada & Dupree 2008; Oklopˇci'c & Hirata 2018), di ff erent relationships have been proposed, aimed at explaining the He i detections, that take into account the F XUV value, either directly (e.g. δ Rp / H eqF XUV; see Fig. 14) or indirectly (e.g. ˙ m obs- ˙ m theory; see Fig 8).  \nHere, we explore other relationships that might explain the He i signals but accounting also for the di ff erences between system characteristics, as δ Rp / H eq. We considered the dimensionless parameter R He / R Hill which is the division between the apparent radius of the planet at the He i triplet wavelength by the Hill radius of the planet. We computed the R He as  \nR He = ( R 2 p + D He R 2 ⋆ ) 1 / 2 , (10)  \nwhere D He is the He i absorption. The equation for the R Hill is  \nR Hill ≃ a p (1 -ecc ) ( M p / (3( M p + M ⋆ ))) 1 / 3 . (11)  \nAs in Eq. 1, we did not take into account planet's eccentricity. Compared to ˙ m obs, R He / R Hill takes into account the He i absorption peak instead of the EW of the signal, and the relation between masses of the system and the a p instead of the star radius. Figure 15 presents R He / R Hill as function of the ˙ m theory. As it is shown, He i detections are aligned in the diagram, hinting that the dimensionless parameter R He / R Hill is a good indicator to investigate He i correlations. R He / R Hill- ˙ m theory have a Pearson correlation r = 0.9034, and we obtained a very similar correlation with ˙ m obs- ˙ m theory (r = 0.8985). For a more quantitative analysis, we fitted the trend that seems to be between R He / R Hill and ˙ m theory. Using the same code as in Sect. 6.3, we fitted the equation R He / R Hill = A log 10 ( ˙ m theory) + log 10 ( B ) and we get A = 0.213 ± 0.027 and B = 2.57 ± 0.15.  \nWe find that gas giant planets are able to fill a larger fraction of their Hill radius (or Hill sphere) than mini-Neptunes. However, He i detections on mini-Neptunes have comparable R He / R Hill values to those from larger planets such as HD189733b or TOI-1268b. This result could point to an observational bias for small planet observations where we are only detecting those exoplanets that their He signal is R He / R Hill ≳ 0.2. Moreover, compared to Fig. 8, some planets move from one side to the other of the line defined by the He i detections, finding planets coded as non-detections above the line or upper limits  \nFig. 15. Relationship between R He / R Hill and the theoretical ( ˙ m theory) energy-limited massloss rate. We define XUV until the He i ionisation range, λ ∈ 5 - 504 Å. The black line indicates the fitted relationship (shown in the legend) and the dashed area the 1 σ uncertainty. He i observations are coded as blue circles for detections, red crosses for nondetections, orange down-pointing triangles for non-conclusive. The colour-coding is according to the results from Sect. 6.3 (see Fig. 8). We did not plot the error bars of ˙ m theory, due to the very large uncertainties associated with the F XUV values and its calculation. Every planet has its name labelled, in black for detections and the rest are in grey.Fig. 16: Histogram of He i triplet observations as function of spectral types (M, K, G, F, and O, B, and A). The number of planets with detections, non-detections, and non-conclusive observations are shown as blue, red, and orange lines, respectively. We indicate the total number of inspected planets in black at the top of each column. We indicate the percentage of detections for each spectral type in blue.  \n<!-- image -->  \nunder the line. TOI-1807 b is a clear example of a non-detection above the line. In particular, the planets with low ˙ m theory value seems to be consistently above the line, which could indicate another bias as well. But, this bias could be more related with the current capabilities of getting more precise measurements. Since they are likely low irradiated by F XUV from the host star, the amount of atoms of He detectable via the NIR triplet might be also very low and very challenging thus for the current telescopes and instruments. The recent He i detection on TOI-2134 b (Zhang et al. 2023a) is consistently in the lower corner on Figures 8 and 15, and might indicate the current limit in terms of signal detectability.", '6.5. He i detections across stellar types': "Oklopˇci'c (2019) indicated that late-type stars are the most favourable host stars to detect the He i triplet on their close-in transiting exoplanets, specially the K-type stars. We explored this theoretical prediction computing the histograms of detections, non-detections and non-conclusive He i triplet observa-  \nons as function of the di ff erent spectral types. Figure 16 shows that K-type exoplanet hosting stars are the most targeted ones (29 planets), but also the spectral type with a larger detection percentage (45%). In fact, the two detections presented in this work, TOI-1268 b and TOI-2018 b, are both around K stars. For M-type stars, GJ 3470 b is the only clear He i triplet detection (Palle et al. 2020a; Ninan et al. 2020), along with the tentative detection on GJ1214b reported by Orell-Miquel et al. 2022. The three detections around F-type stars (HD 209458 b, Alonso-Floriano et al. 2019; HAT-P-67 b, Gully-Santiago et al. 2024; and HAT-P-32 b, Czesla et al. 2022) are interesting because these planets are hot Jupiters, and HAT-P-67 b and HAT-P-32 b possess very extended He structures (Gully-Santiago et al. 2024; Zhang et al. 2023c). To date, there are no He i detections of exoplanets orbiting G, A, B, or O stars. To further understand the dependence on spectral type, we also computed the histogram of He i detections and non-detections as function of the stellar mass (Fig. A.23). The number of detections increases from 0.47 M ⊙ (GJ 3470) until peaking at 0.9 M ⊙ . Above 0.9 M ⊙ , non-detections are dominant over detections. Therefore, we can estimate a limit in host stellar mass at ∼ 0.9 M ⊙ for the He i detections, which deserves further investigation.  \nWe note that all the He i detections from Table M.1 are found in planets orbiting stars with T e ff ≲ 6250 K, which is the temperature at which the Kraft Break happens (Kraft 1967). Exoplanet spin-orbit alignment seems to be related to the Kraft Break (e.g. Winn et al. 2010; Brown et al. 2017; Attia et al. 2023). To date, the hottest host star with He i detection is HATP-67 with T e ff ∼ 6400 K. The observations of other ten planets orbiting ≳ 6250 K stars resulted in upper limits only. The Kraft Break marks also the transition between stars with outer convection zones and those without. These convection zones are predicted to generate magnetic fields able to heat the chromosphere and the corona, increasing the XUV radiation (Wright et al. 2011). Therefore, stars without this heating mechanism are less likely to significantly populate the He metastable level of their exoplanets' atmospheres. The link between the stellar T e ff , the flux in the XUV range and the He i triplet detections deserves further investigation.", '6.6. He i dependence on planetary parameters': "One of the aims of the MOPYS project was also to find a relation between planetary and stellar properties that could explain, and predict, the He i detections and non-detections, but without the need to involve the strength of the signal itself. From the many explored relations and diagrams, only the equilibrium temperature ( T eq) and the planet density ( ρ p) showed some hints of such relationship. Figure 17 presents the T eq versus ρ p diagram where the information about the He i observations is encoded in the colour markers (blue circles are detections, red crosses are non-detections, and non-conclusive observations are not shown). Although both are planet properties, T eq encapsulates information about the star via the T e ff and the a p / R ⋆ parameters.  \nFrom Fig. 17, one clear result is that there are no He i detections on exoplanets with T eq ≳ 2000 K. The hottest detection at T eq ∼ 2000 K comes from the pu ff y planet HAT-P-67 b. In fact, the data in Figure 17 hint to a density / T eq upper boundary for the detections. We draw 'by-hand' this observational limit for a better visualisation, which follows T eq ≃-500 log10( ρ p) trend, approximately. If we assume that detections and non-detections are distributed randomly across the T eqρ p parameter space, the probability of extracting 14 non-detections from a pool of 17 detections and 36 non-detections is only 0.16 %. Although it is an  \nFig. 17: Equilibrium temperature ( T eq) vs planet density ( ρ p) diagram for the He i detections (blue circles) and non-detections (red crosses). The black line (drawn by-eye) indicates the hinted upper boundary for the He i detections. The coloured dashed vertical lines indicate the Earth's density ( ρ ⊕ ; orange line at ∼ 5.5 g cm -3 ), and a representative density where the pu ff y planet population starts ( ρ pu ff y; blue line at ∼ 1.3 g cm -3 ; Luque & Pallé 2022). All known planets with T eq and ρ p determined with a precision better than 50% are marked with grey dots (data from NASA Exoplanet Archive).  \n<!-- image -->  \n1  \napproximate calculation, the very low probability supports the robustness of the boundary. For planets falling above the line it is unlikely that one can detect He i , as all the inspected planets resulted in non-detections. Below that line, He i detections and non-detections are equally found.  \nAt face value our results indicate that for high-density planets (from Earth-size to sub-Neptune-size) He i absorptions are present only in cool atmospheres. A possible interpretation is that these type of planets quickly lose any H / He atmosphere when subjected to strong stellar instellation (assuming that T eq is directly related at the population level to stellar flux). For low-density planets (gas giants) He i detection are found up to ∼ 2000 K.  \nFurther He i detections will confirm or disprove the observational boundary proposed here, and define its exact dependence on di ff erent planet parameters.", '7. Conclusions': "This work shows the first results of the MOPYS survey, a project that aims to test the predictions of atmospheric evolution theories by confronting them with observations of evaporation proxies. We present CARMENES and GIARPS high-resolution transmission spectroscopy observations of 20 exoplanets. We analysed their atmospheres searching for He i triplet absorption in the NIR, and H α in the VIS as well. We report two new detections of the He i triplet for TOI-1268 b and TOI-2018 b, and a new H α detection for TOI-1136 d. We also found hints of He i on HD63433b, and of H α on HD73583b and c, which need further confirmation.  \nWe complemented our target list with other planets from the literature with He i triplet and H α (and Ly α for some particular planets as well) observations, for a total of 70 planets. We considered those lines as proxy of evaporation to test mass-loss theories. The main findings of our study are:  \n- -Our age distribution of evaporation detections does not favour either photo-evaporation or core-powered mass-loss as planet formation mechanisms.\n- -We find no trend in atmospheric evaporation with stellar age. Young ( < 1-Gyr-old) planets do not exhibit more He i or H α detections than older planets.\n- -Evaporation (He i or H α ) signals are more frequent or easier to detect for planets around stellar hosts within ∼ 1-3 Gyr ages.\n- -The fraction of planets that show He i detections is much larger for K-type stars (45%) than for any other spectral type.\n- -We find no evaporation detections of young and old planets below the radius gap, confirming rocky planets are not under extreme evaporation, at least after the first ∼ 300 Myr.\n- -We find hints that evaporation of pu ff y sub-Neptunes ( R p > 3 R ⊕ ) happens at ages older than 1 Gyr, although the number of planets is small, and a larger sample is needed to confirm this finding.\n- -Weprovide new constraints to the cosmic shoreline, by using the evaporation detections as evidences for the existence of planetary atmospheres. Our cosmic shoreline move to higher irradiation flux, reducing the parameter space of bare rocky planets.\n- -We present the He-related dimensionless parameter R He / R Hill, as a new valid parameter to study the He i detections and upper limits.\n- -All the He i detections are found in planets orbiting stars with T e ff ≲ 6250 K, which is the temperature at which the Kraft Break happens.\n- -We determine a statistically significant observational upper boundary for He i detections in the T eq versus ρ p parameter space. The line decays at a rate of T eq ≃-500 log10( ρ p), approximately. Planets falling above that boundary are unlikely to show He i absorption signals.  \nWe encourage further evaporation observations (Ly α , H α , and He i triplet) to increase and complete the sample presented in this work. Although this work represents the biggest evaporation survey to date (specially focused on < 1Gyr-old planets), some key questions of planet formation and atmospheric evaporation remained unanswered. In particular, the detection and atmospheric characterisation of very young planets ( < 100 Myr) is key for discriminating between di ff erent mass-loss processes in planetary atmospheres. It is precisely because of their age that they are the most di ffi cult planets to analyse their atmospheres.  \nAcknowledgements. CARMENES is an instrument at the Centro Astronómico Hispano en Andalucía (CAHA) at Calar Alto (Almería, Spain), operated jointly by the Junta de Andalucía and the Instituto de Astrofísica de Andalucía (CSIC). The authors wish to express their sincere thanks to all members of the Calar Alto sta ff for their expert support of the instrument and telescope operation. CARMENES was funded by the Max-Planck-Gesellschaft (MPG), the Consejo Superior de Investigaciones Científicas (CSIC), the Ministerio de Economía y Competitividad (MINECO) and the European Regional Development Fund (ERDF) through projects FICTS-2011-02, ICTS-2017-07-CAHA4, and CAHA16-CE-3978, and the members of the CARMENES Consortium (Max-Planck-Institut für Astronomie, Instituto de Astrofísica de Andalucía, Landessternwarte Königstuhl, Institut de Ciències de l'Espai, Institut für Astrophysik Göttingen, Universidad Complutense de Madrid, Thüringer Landessternwarte Tautenburg, Instituto de Astrofísica de Canarias, Hamburger Sternwarte, Centro de Astrobiología and Centro Astronómico Hispano-Alemán), with additional contributions by the MINECO, the Deutsche Forschungsgemeinschaft  \n(DFG) through the Major Research Instrumentation Programme and Research Unit FOR2544 'Blue Planets around Red Stars', the Klaus Tschira Stiftung, the states of Baden-Württemberg and Niedersachsen, and by the Junta de Andalucía. Based on observations collected at the Centro Astronómico HispanoAlemán (CAHA) at Calar Alto, proposals 21B-3.5-004, 22A-3.5-007, 22B-3.5009, 23A-3.5-009, and 23B-3.5-002, operated jointly by Junta de Andalucía and Consejo Superior de Investigaciones Científicas (IAA-CSIC). We used data from the CARMENES data archive at CAB (CSIC-INTA). We acknowledge the use of the ExoAtmospheres database during the preparation of this work This work is partly supported by JSPS KAKENHI Grant Number JPJP24H00017 and JSPS Bilateral Program Number JPJSBP120249910. This paper is based on observations made with the MuSCAT2 instrument, developed by ABC, at Telescopio Carlos Sánchez operated on the island of Tenerife by the IAC in the Spanish Observatorio del Teide. We acknowledge financial support from the Agencia Estatal de Investigación (AEI / 10.13039 / 501100011033) of the Ministerio de Ciencia e Innovación and the ERDF 'A way of making Europe' through projects PID2022-137241NBC4[1:4], PID2019-109522GB-C5[1:4], PID2021125627OB-C31, PID2022-141216NB-I00, Ariel Postdoctoral Fellowship program of the Swedish National Space Agency (SNSA), and the Centre of Excellence 'Severo Ochoa' and 'María de Maeztu' awards to the Instituto de Astrofísica de Canarias (CEX2019-000920-S), Instituto de Astrofísica de Andalucía (CEX2021-001131-S) and Institut de Ciències de l'Espai (CEX2020001058-M). J.O-M agraeix el recolzament, suport i ànims que sempre ha rebut per part de Padrina Conxa, Padrina Mercè, Jeroni, Mercè i més familiars i amics. J.O.M. gratefully acknowledge the inspiring discussions with Maite Mateu, Alejandro Almodóvar, and Joan Perelló, and the support from Guillem, Benet, and Montse. This research has made use of resources from VallAlbaida-Mallorca collaboration. J.O.M. acknowledges the contributions and patience of Jorge Terol Calvo, i molt especialment a tu, Yess. L.N. and F.L. acknowledge the support by the Deutsche Forschungsgemeinschaft (DFG, German Research Foundation) Project number 314665159. E.K. and S.H.A. acknowledge the support from the Danish Council for Independent Research through grant No.2032-00230B. E.N. acknowledges the support by the DFG Research Unit FOR2544 'Blue Planets around Red Stars'. S.C. acknowledges the support of the DFG priority program SPP 1992 'Exploring the Diversity of Extrasolar Planets' (CZ 222 / 5-1).", 'References': "Addison, B. C., Knudstrup, E., Wong, I., et al. 2021, AJ, 162, 292  \n- Allan, A. P., Vidotto, A. A., Villarreal D'Angelo, C., Dos Santos, L. A., & Driessen, F. 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V., Gully-Santiago, M., et al. 2023c, Science Advances, 9, eadf8736  \nZicher, N., Barragán, O., Klein, B., et al. 2022, MNRAS, 512, 3060  \n- 1 Instituto de Astrofísica de Canarias (IAC), 38205 La Laguna, Tenerife, Spain\n- e-mail: jom@iac.es\n- 2 Departamento de Astrofísica, Universidad de La Laguna (ULL), 38206 La Laguna, Tenerife, Spain\n- 3 Instituto de Astrofísica de Andalucía (IAA-CSIC), Glorieta de la Astronomía s / n, E-18008 Granada, Spain\n- 4 Centro de Astrobiología (CSIC-INTA), Camino Bajo del Castillo s / n, Villanueva de la Cañada, E-28692 Madrid, Spain\n- 5 Institut für Astrophysik und Geophysik, Georg-August-Universität, Friedrich-Hund-Platz 1, 37077 Göttingen, Germany\n- 6 Thüringer Landessternwarte Tautenburg, Sternwarte 5, 07778 Tautenburg, Germany\n- 7 Institut de Ciències de l'Espai (ICE, CSIC), Campus UAB, Can Magrans s / n, 08193 Bellaterra, Barcelona, Spain\n- 8 Institut d'Estudis Espacials de Catalunya (IEEC), 08034 Barcelona, Spain\n- 9 Osservatorio Astronomico di Padova, Vicolo dell'Osservatorio 5, 35122, Padova, Italy\n- 10 Astrobiology Center, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan\n- 11 National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan\n- 12 Astronomical Science Program, Graduate University for Advanced Studies, SOKENDAI, 2-21-1, Osawa, Mitaka, Tokyo, 181-8588, Japan\n- 13 Department of Space, Earth and Environment, Chalmers University of Technology, 412 93, Gothenburg, Sweden\n- 14 Stellar Astrophysics Centre, Department of Physics and Astronomy, Aarhus University, Ny Munkegade 120, DK-8000 Aarhus C, Denmark\n- 15 NASA Sagan Fellow\n- 16 Division of Geological and Planetary Sciences, 1200 E California Blvd, Pasadena, CA, 91125, USA\n- 17 Department of Astronomy, California Institute of Technology, Pasadena, CA 91125, USA\n- 18 Komaba Institute for Science, The University of Tokyo, 3-8-1 Komaba, Meguro, Tokyo 153-8902, Japan\n- 19 Faculty of Physics, Ludwig Maximilian University, Scheinerstrasse 1, D-81679, Munich, Bavaria, Germany\n- 20 Max-Planck-Institute für Astronomie, Königstuhl 17, D-69117 Heidelberg, Germany\n- 21 Department of Multi-Disciplinary Sciences, Graduate School of Arts and Sciences, The University of Tokyo, 3-8-1 Komaba, Meguro, Tokyo\n- 22 Departamento de Física de la Tierra y Astrofísica and IPARCOSUCM (Instituto de Física de Partículas y del Cosmos de la UCM), Facultad de Ciencias Físicas, Universidad Complutense de Madrid, E-28040, Madrid, Spain\n- 23 Landessternwarte, Zentrum für Astronomie der Universität Heidelberg, Königstuhl 12, D-69117 Heidelberg, Germany  \n- 24 Hamburger Sternwarte, Gojenbergsweg 112, 21029 Hamburg, Germany\n- 25 Centro Astronómico Hispano en Andalucía, Observatorio Astronómico de Calar Alto, Sierra de los Filabres, 04550 Gérgal, Almería, Spain", 'Appendix B: V1298Tau extra material': 'Table B.1: Prior and posterior distributions from the nested sampling fitting for V1298 Tau c H α feature (see Fig. A.2). Prior label U represents uniform distribution.Table B.2: Prior and posterior distributions from the nested sampling fitting for V1298 Tau c He i feature (see Fig. A.2). Prior label U represents uniform distribution.  \nFig. A.1: Same as Fig. 1, but for MASCARA-2 b observations with CARMENES.  \n<!-- image -->  \nFig. A.2: Same as Fig. 1, but for V1298 Tau c observations with CARMENES.  \n<!-- image -->  \nFig. A.3: Same as Fig. 1, but for TOI-1431 b observations with CARMENES.  \n<!-- image -->  \nFig. A.4: Same as Fig. 1, but for TOI-2048 b observations with CARMENES.  \n<!-- image -->  \nFig. A.5: Same as Fig. 1, but for HD 63433 b observations with CARMENES.  \n<!-- image -->  \nFig. A.8: Same as Fig. 1, but for HD 73583 c observations with GIARPS.  \n<!-- image -->  \nFig. A.9: Same as Fig. 1, but for HD 235088 b observations with CARMENES.  \n<!-- image -->  \nFig. A.10: Same as Fig. 1, but for K2-77 b observations with CARMENES.  \n<!-- image -->  \nFig. A.11: Same as Fig. 1, but for TOI-2046 b observations with GIARPS.  \n<!-- image -->  \nFig. A.12: Same as Fig. 1, but for TOI-1807 b observations with CARMENES.  \n<!-- image -->  \nFig. A.13: Same as Fig. 1 for TOI-1136 d H α observations with HARPS-N (14-May-2021) and CARMENES VIS (30-Jan-2023).  \n<!-- image -->  \nFig. A.14: Same as Fig. 1, but for TOI-2076 b observations with CARMENES.  \n<!-- image -->  \nFig. A.15: Same as Fig. 1, but for TOI-1683 b observations with GIANO-B.  \n<!-- image -->  \nFig. A.16: Same as Fig. 1, but for WASP-189 b observations with GIANO-B.  \n<!-- image -->  \nFig. A.17: Same as Fig. 1, but for HAT-P-57 b observations with CARMENES.  \n<!-- image -->  \n3  \n3.00  \nFig. A.19: Same as Fig. 1, but for TOI-2018 b idividual observations of He i triplet with GIANO-B.  \n<!-- image -->  \nFig. A.20: Time evolution of the activity indicators for K2-100 b (top left), V1298 Tau c (top right), TOI-2048 b (middle left), HD63433b (1 November 2021, middle right; and 27 November 2022, bottom left), and HD 63433 c (bottom right). The vertical dashed and dotted lines represent the di ff erent contacts during the transit.  \n<!-- image -->  \nFig. A.21: Same as Fig. A.20, but for HD 73583 b (top left), HD 73583 b (top right), K2-77 b (middle left), TOI-2046 b (middle right), and TOI-1807 b (16 December 2021, bottom left; and 23 December 2022, bottom right). Article number, page 40 of 64  \n<!-- image -->  \n100  \nFig. A.22: Same as Fig. A.20, but for TOI-1136 d (14 May 2021, top left; and 30 January 2023, top right), TOI-1268 b (middle left), and TOI-2076 b (middle right), and TOI-2018 b (9 April 2022, bottom left; and 15 June 2022, bottom right).  \n<!-- image -->  \nArticle number, page 41 of 64  \nFig. A.23: Histogram of He i triplet observations (detections: blue, non-detections: red) as a function of stellar mass.  \n<!-- image -->  \nFig. B.1: Corner plot for the nested sampling posterior distribution of V1298 Tau c H α (left) and He i (right) feature.  \n<!-- image -->  \nFig. B.2: TS around the H α (left) and He i triplet (right) lines for the transit phases from top to bottom (consecutively o ff set): start (blue), centre (red), end (green), and post-transit (orange). Master-Out spectrum is overplotted in black along with the individual TS for comparison. The dotted vertical lines indicate the lines positions. The wavelengths in this figure are given in vacuum.  \n<!-- image -->', 'Appendix C: TOI-2048 extra material': "We analysed the TESS data from Sectors 16, 23, 24, 50, and 51 with the procedure described in Sect. 2.2. We used the celerite GPquasi-periodic kernel to account for the young star variability and we adopted the stellar parameters from Newton et al. (2022) for a proper comparison.  \nThe fitted parameters with their prior and posterior values, and the derived parameters for TOI-2048 b are shown in Table C.1. The TESS data along with the best transiting and GP models is shown in Fig. C.1 and TOI-2048 b phase folded transit is shown in Fig. C.2. We obtained similar results using the celerite GP exponential kernel.  \nTo derive an estimation of TOI-2048 b's semi-amplitude K ⋆ , first we forecasted its mass using the probabilistic mass-radius relationship for sub-Neptune-sized planets ( R p < 4 R ⊕ ) of Wolfgang et al. (2016). We predicted a planetary mass of ∼ 9 ± 2 M ⊕ , which is translated into K ⋆ ≃ 2.8 ± 0.6 m s -1 using the equation  \nK ⋆ = 28 . 4 m s -1 ( P pl / year) -1 / 3 ( M pl / MJup) ( M ⋆/ M ⊙ ) -2 / 3 (C.1)  \nTable C.1: Prior and posterior distributions from the juliet fitting for TOI-2048 b. Prior labels U , N , F , and J represents uniform, normal, fixed, and Je ff rey's distribution, respectively.", 'Derived planetary parameters': "Notes. ( a ) Central time of transit ( t 0) units are BJD -2 457 000. ( b ) Equilibrium temperatures were calculated assuming zero Bond albedo.  \nWe analysed TOI-1683 TESS data from Sectors 19, 43, and 44. Because our purpose is to model the stellar rotation, we used the celerite GP quasi-periodic kernel. We adopted the stellar parameters used in Zhang et al. (2023b) for a proper comparison. The fitted parameters with their prior and posterior values, and the derived parameters for the planet candidate TOI-1683 b are shown in Table K.1. The TESS data along with the best transiting and GP models is shown in Fig. K.2 and TOI-1683 b phase folded transit is shown in Fig. K.1.  \nWe computed the GLS periodogram (Zechmeister & Kürster 2009, generalised Lomb-Scargle) in the light curve of the three available TESS Sectors. We derived a stellar rotation period of 22.5 ± 5.2 days, which is consistent with the GPP rot value from the photometric fit (GPP rot 20 ± 1 d). We estimated TOI-1683's age using the gyrochronology methods from Mamajek & Hillenbrand (2008, Eq. 12, 13, 14 and parameters from Table 10), Schlaufman (2010, Eq. 1), and Bouma et al. (2023, python package gyro-interp 8 ), and the derived ages are 1450 + 700 -550 Myr, 2300 + 1300 -900 Myr, and 2600 ± 1200 Myr, respectively. All the three methods are consistent with an age older than 1 Gyr. Figure K.3  \nFig. K.1: TOI-1683 TESS photometry (blue dots with error bars) phase-folded to the period P and central time of transit t 0 (shown above the panel, t 0 units are BJD -2 457 000) derived from the juliet fit. The black line is the best transit model for TOI1683 b. The orange points show binned photometry for visualisation.  \n<!-- image -->  \nPb =3.058 d, t 0 = 2522.700  \nPhase  \nResiduals (ppm)  \nshows the distribution of rotation periods as a function of the G -J colour for various young clusters. From this qualitative analysis, it is evident that TOI-1683 appears to be older than the NGC6811cluster (which has an age ∼ 1 Gyr), and it is positioned above the sequence of NGC 6774 ( ∼ 2.5 Gyr).  \nOn the other hand, the presence of atmospheric absorption lines, such as Li i at 6709.61 Å, can serve as a valuable age indicator for the stars. We searched for the presence of Li i in the coadded spectrum generated by the serval from the CARMENES spectra. However, no clear Li i feature was detected in the spectrum. Therefore, we set an upper limit for the Li i equivalent width of 1 mÅ at 3 σ . As Figure K.4 shows, this upper limit suggests that the star is older than the Hyades or Praesepe clusters, which have ages between 590-650 Myr. Lastly, to investigate the kinematic properties of TOI-1683 and to determine if it shares any characteristics with known clusters, moving groups, or associations, we calculated the UVW galactocentric space velocities using the astrometry and systemic velocity data from Gaia . Figure K.5 places TOI-1683 outside of any shown young moving group or cluster. Additionally, we conducted a search in the literature of cluster catalogs, and TOI-1683 was not found to be associated with any of these known stellar groups. These results suggest that it is not part of any well-established stellar association or cluster. Based on our analysis, we adopt an age older than 1 Gyr for TOI-1683 and, therefore, it is not a young object in terms of planet formation and evolution.  \nThe age resulting of combining the three gyrochronology methods is 2000 + 1300 -900 Myr, which is the age used in this work. However, Zhang et al. (2023b) claim that TOI-1683 is a young star ( < 1 Gyr) with an age of 500 ± 150 Myr based on its gyrochronology analysis. The main di ff erence with our analyses  \nFig. K.2: TOI-1683 2 min cadence TESS photometry from Sectors 19, 43, and 44 along with the transit plus GP model. The upward-pointing red triangles mark the transit times for the TOI-1683 b.  \n<!-- image -->  \nFig. K.3: Rotation period distribution as a function of colour G -J for the Pleiades ( ∼ 125 Myr; Rebull et al. 2016), Praesepe ( ∼ 590 Myr; Douglas et al. 2017), Hyades ( ∼ 650 Myr; Douglas et al. 2019), NGC 6811 ( ∼ 1000 Myr; Curtis et al. 2019), and NGC 6774 ( ∼ 2500 Myr; Gruner & Barnes 2020) clusters. The gold star represents TOI-1683.  \n<!-- image -->  \n-  \nis that they derived a rotation period of 11.3 ± 1.5 days, which is consistent with the half of the rotation period that we derived. Looking at a single TESS sector it is possible to confuse the stellar rotation period with one of its harmonics when it is of the order of the duration of the TESS sector ( ∼ 28 days). This could be the case of TOI-1683 Sector 19. However, the consecutive Sectors 43 and 44 make it clear that the stellar rotation period of TOI-1683 is greater than 11.3 days.  \nFig. K.4: Equivalent width distribution of Li i as a function of the e ff ective temperature for the Pleiades ( ∼ 125 Myr; Bouvier et al. 2018), Praesepe ( ∼ 590 Myr), and Hyades ( ∼ 650 Myr; Cummings et al. 2017). The gold triangle represents TOI-1683.  \n<!-- image -->  \nFig. K.5: UVW velocity diagram for TOI-1683 (gold star). The members of the Castor moving group, Hyades supercluster (Hs), IC 2391 supercluster, Local Association (LA), and Ursa Major group (UMa) from Montes et al. (2001) are included. The ellipses represent the 3 σ values of the UVW for each Young Moving Group.  \n<!-- image -->  \nTable L.1: Prior and posterior distributions from the nested sampling fitting for TOI-2018 b He i signal from the first night. Prior label U represents uniform distribution.Table L.2: Prior and posterior distributions from the nested sampling fitting for TOI-2018 b He i signal from the second night. Prior label U represents uniform distribution.", 'Appendix D: HD63433 system extra material': 'Fig. D.1: Contribution of the RM and CLV e ff ects to the HD 63433 b transmission spectra around the H α (top) and He i NIR triplet (bottom) lines. The vertical cyan lines mark the position of the lines of interest.  \n<!-- image -->  \nTable D.1: Prior and posterior distributions from the nested sampling fitting for HD 63433 b H α feature (see Fig. A.5). Prior label U represents uniform distribution.  \nFig. D.2: Corner plot for the nested sampling posterior distribution of HD 63433 b H α feature from partial transit.  \n<!-- image -->', 'Appendix E: HD73583 system extra material': "We analysed HD73583 TESS data from Sectors 8, 34, and 61. Because our purpose for this system is not to study the young stellar activity but derive precise ephemeris, we used the celerite GP exponential kernel to account for the stellar variability and rotation. We adopted the stellar parameters used in Barragán et al. (2022).  \nThe fitted parameters with their prior and posterior values, and the derived parameters for HD 73583 b and c are shown in Table E.1. The TESS data along with the best transiting and GP models is shown in Fig. E.2 and HD 73583 b and c phase folded transits are shown in Fig. E.1.  \nTable E.1: Prior and posterior distributions from the juliet fitting for HD 73583 b and c. Prior labels U , N , F , and J represents uniform, normal, fixed, and Je ff rey's distribution, respectively.  \nNotes. ( a ) Central time of transit ( t 0) units are BJD -2 457 000. ( b ) Equilibrium temperatures were calculated assuming zero Bond albedo.  \nResiduals (ppm)  \nPb  \n=6.398 d,  \nt  \n0 = 2592.563  \nPc  \n=18.879 d,  \nt  \n0 = 2949.582  \nFig. E.1: HD 73583 TESS photometry (blue dots with error bars) phase-folded to the period P and central time of transit t 0 (shown above each panel, t 0 units are BJD -2 457 000) derived for planet b (left), and c (right) from the juliet fit. The black line is the best transit model for each planet. The orange points show binned photometry for visualisation.  \n<!-- image -->  \nFig. E.2: HD 73583 2 min cadence TESS photometry from Sectors 8, 34, and 61 along with the transit plus GP model. The upwardpointing red and green triangles mark the transit times for HD 73583 b and c, respectively.  \n<!-- image -->  \nTable E.2: Prior and posterior distributions from the nested sampling fitting for HD 73583 b H α feature (see Fig. A.7). Prior label U represents uniform distribution.  \nTable E.3: Prior and posterior distributions from the nested sampling fitting for HD 73583 c H α feature (see Fig. A.8). Prior label U represents uniform distribution.  \nFig. E.3: Corner plot for the nested sampling posterior distribution of HD 73583 b (left) and HD 73583 c (right) H α features.  \n<!-- image -->", 'Appendix F: K2-77b planetary parameters': "We analysed the K2-77 b TESS data from Sectors 5, 42, 43, and 44. Because we did not see a clear rotation modulation in the light curves, we used the celerite GP exponential kernel to account for the stellar variability. We adopted the stellar parameters used in Gaidos et al. (2017).  \nThe fitted parameters with their prior and posterior values, and the derived parameters for K2-77 b are shown in Table F.1. The TESS data along with the best transiting and GP models is shown in Fig. F.2 and K2-77 b phase folded transit is shown in Fig. F.1.  \nAs we did in App. C, we forecasted K2-77 b mass and we computed the estimated K ⋆ using Eq. C.1. We predicted a planetary mass of ∼ 9 ± 2 M ⊕ , and a semi-amplitude K ⋆ of ∼ 3.4 ± 0.8 m s -1 . Gaidos et al. (2017) only reported an upper limit to K2-77 b's mass of 1.9 MJ .  \nTable F.1: Prior and posterior distributions from the juliet fitting for K2-77 b. Prior labels U , N , F , and J represents uniform, normal, fixed, and Je ff rey's distribution, respectively.", 'Pb =8.200 d, t 0 = 2522.634': 'Phase  \nResiduals (ppm)  \n1.02  \nFig. F.2: K2-77 2 min cadence TESS photometry from Sectors 5, 42, 43, and 44 along with the transit plus GP model. The upwardpointing red triangles mark the transits for K2-77 b.  \n<!-- image -->', 'Appendix G: TOI-1807b extra material': "We analysed TOI-1807 TESS data from Sectors 22, 23, and 49. Because our purpose is not to model the young stellar activity, but derive precise ephemerides, we used the unspecific celerite GP exponential kernel to account for the stellar variability and rotation. We adopted the stellar parameters used in Nardiello et al. (2022) for a proper comparison.  \nThe fitted parameters with their prior and posterior values, and the derived parameters for TOI-1807 b are shown in Table G.1. The TESS data along with the best transiting and GP models is shown in Fig. G.2 and TOI-1807 b phase folded transit is shown in Fig. G.1.  \nTable G.1: Prior and posterior distributions from the juliet fitting for TOI-1807 b. Prior labels U , N , F , and J represents uniform, normal, fixed, and Je ff rey's distribution, respectively.", 'Appendix H: TOI-1136d extra material': 'Fig. H.1: Contribution of the RM and CLV e ff ects to the TOI1136 d transmission spectra around the H α (top) and He i NIR triplet (bottom) lines. The vertical cyan lines mark the position of the lines of interest.  \n<!-- image -->  \nTable H.1: Prior and posterior distributions from the nested sampling fitting for TOI-1136 d H α feature (see Fig. 2). Prior label U represents uniform distribution.  \nFig. H.2: Corner plot for the nested sampling posterior distribution of TOI-1136 d H α feature.  \n<!-- image -->', 'Appendix I: TOI-1268 b extra material': 'Fig. I.1: MuSCAT2 ground-based transit observation of TOI1268 b. Raw light curve (left panel), detrended light curve (centre panel), and residuals (right panel) for the Sloan g (blue), r (green), i (yellow), zs (red) filters. The points show the individual observations, the circles represent binned data points, and the best-fit model is shown by the black line.  \n<!-- image -->  \nTable I.1: Prior and posterior distributions from the nested sampling fitting for TOI-1268 b He i signal (see Fig. 4). Prior label U represents uniform distribution.  \nFig. I.2: Corner plot for the nested sampling posterior distribution of TOI-1268 b He i signal.  \n<!-- image -->  \nResiduals (ppm)', 'Appendix J: TOI-2076b extra material': "Table J.1: Prior and posterior distributions from the juliet fitting for TOI-2076 b data from LCO (see Fig. J.1). Prior labels U , N , F , and J represents uniform, normal, fixed, and Je ff rey's distribution, respectively.  \nNotes. ( a ) Central time of transit ( t 0) units are BJD -2 457 000. ( b ) Equilibrium temperatures were calculated assuming zero Bond albedo.  \nPb  \n=10.355 d,  \nt  \n0 = 2460079.549  \nFig. J.1: TOI-2076 LCO photometric data (blue dots with error bars). The period P and central time of transit t 0 derived from the juliet fit are shown above the panel. The black line is the best transit model for TOI-2076 b. The orange points show binned photometry for visualisation. The GP contribution was removed from the data.  \n<!-- image -->", 'Appendix K: TOI-1683b extra material': "Table K.1: Prior and posterior distributions from the juliet fitting for TOI-1683 b. Prior labels U , N , F , and J represents uniform, normal, fixed, and Je ff rey's distribution, respectively.", 'Appendix L: TOI-2018b extra material': 'Fig. L.1: MuSCAT2 ground-based transit observation of TOI2018 b. Raw light curve (left panel), detrended light curve (centre panel), and residuals (right panel) for the Sloan g (blue), r (green), i (yellow), zs (red) filters. The points show the individual observations, the circles represent binned data points, and the best-fit model is shown by the black line.  \n<!-- image -->  \nFig. L.2: Corner plot for the nested sampling posterior distribution of TOI-2018 b He i signals for first (top) and second (bottom) nights.  \n<!-- image -->  \nFig. L.3: Corner plot for the nested sampling posterior distribution of TOI-2018 b He i signals from the combined nights.  \n<!-- image -->  \nTable L.3: Prior and posterior distributions from the nested sampling fitting for TOI-2018 b He i signal from the combined nights. Prior label U represents uniform distribution.', 'Appendix M: He i database information': 'As Zhang et al. (2023a) remarked, the compilation of consistent parameters and r tudy the He i triplet is a challenging work. We thank the e ff ort done by Zhang et al. (2023a) to construct Table 3 therein. We based our He i database (Table M.1) on the compilations by Kirk et al. (2022), Czesla et al. (2022), Fossati et al. (2022, 2023), Guilluy et al. (2023), Zhang et al. (2023a), and J. Sanz-Forcada, priv. comm., which we further complemented.  \nBecause the He i results are given in terms of the equivalent width (EW) or absorption peak of the line, we firstly needed a relation between both magnitudes to convert one to the other. The EW for a spectral line is the width of a rectangle, with a height of the continuum level, such that the area of the spectral line is equal to the area of the rectangle. Usually, the planetary spectral lines are fitted using a Gaussian profile as G( λ ) = A exp ( -( λ -λ 0) 2 / (2 σ 2 )) where A is the amplitude of the signal, λ 0 is the central position of the line, and σ is the standard deviation. Thus, the area under the curve described by G( λ ) can be computed as  \n<!-- image -->  \nIf the spectra are normalised to 1, we can calculate EW = | A | √ 2 πσ 2 , where A is in units of normalised flux. EW is in the units of σ . Then, we calculated from the papers where EWand absorption values are reported (e.g. Guilluy et al. 2020; Casasayas-Barris et al. 2021a; Zhang et al. 2022b, 2023b; OrellMiquel et al. 2022, 2023) a masterσ as the mean value. We obtained masterσ = 0.43 mÅ.', 'T able M.1: continued.': "T able M.1: continued.  \nNotes. a Y : Detection, ND: Non Detection, NC: Non Conclusi v e. F or H, we indicate if the detection / non-detection comes from H α (H), L y α (L), or both (B). b F XUV from: 1 Sanz-F orcada in prep., 2 P alle et al. ( 2020a ), 3 Zhang et al. ( 2023a ), 4 Guilluy et al. ( 2023 ), 5 Nortmann et al. ( 2018 ), 6 Alonso-Floriano et al. ( 2019 ), 7 Orell-Miquel et al. ( 2023 ), 8 Allart et al. ( 2018 ), 9 V issaprag ada et al. ( 2022b ), 10 Czesla et al. ( 2022 ), 11 Allart et al. ( 2023 ), 12 Czesla et al. ( 2024 ), 13 Kirk et al. ( 2022 ), 14 dos Santos et al. ( 2020 ), and 15 V issaprag ada et al. ( 2024b ) Refer ences. Data compilation from N ASA Exoplanet Archi v e and ExoAtmospheres database. Some references in particular are as follo ws. 55 Cnc e: v on Braun et al. ( 2011 ); Zhang et al. ( 2021 ) ; A U Mic b: Pla vchan et al. ( 2020 ); Zicher et al. ( 2022 ); P alle et al. ( 2020b ); Hirano et al. ( 2020 ); Rockcli ff e et al. ( 2023 ) ; GJ436 b: Nortmann et al. ( 2018 ); Caule y et al. ( 2017 ); K ulo w et al. ( 2014 ); Ehrenreich et al. ( 2015 ) ; GJ806 b: P alle et al. ( 2023 ) ; GJ1214 b: Kasper et al. ( 2020 ); Orell-Miquel et al. ( 2022 ); Spak e et al. ( 2022 ) ; GJ3470 b: Bourrier et al. ( 2018 ); P alle et al. ( 2020a ); Ninan et al. ( 2020 ) ; GJ9827 b: Krishnamurth y et al. ( 2023 ); Carleo et al. ( 2021 ) ; GJ9827 d: Krishnamurth y et al. ( 2023 ); Carleo et al. ( 2021 ); Kasper et al. ( 2020 ) ; HD63433 b: Mann et al. ( 2020 ); Mallorquín et al. ( 2023 ); Zhang et al. ( 2022c ) ; HD63433 c: Mann et al. ( 2020 ); Mallorquín et al. ( 2023 ); Zhang et al. ( 2022c ) ; HD73583 b: Barragán et al. ( 2022 ); Zhang et al. ( 2022b , 2023b ) ; HD73583 c: Barragán et al. ( 2022 ) ; HD89345 b: Guilluy et al. ( 2023 ) ; HD97658 b: Kasper et al. ( 2020 ) ; HD189733 b: Jensen et al. ( 2012 ); Salz et al. ( 2018 ); Guilluy et al. ( 2020 ); Zhang et al. ( 2022a ); Allart et al. ( 2023 ); Dos Santos et al. ( 2023 ) ; HD209458 b: V idal-Madjar et al. ( 2003 ); Alonso-Floriano et al. ( 2019 ) ; HD235088 b: Orell-Miquel et al. ( 2023 ) ; HA T -P-3 b: Guilluy et al. ( 2023 ) ; HA T -P-11 b: Allart et al. ( 2018 ); Ben-Ja ff el et al. ( 2022 ); Allart et al. ( 2023 ) ; HA T -P-18 b: Fu et al. ( 2022 ); F ournier -T ondreau et al. ( 2024 ) ; HA T -P-32 b: Czesla et al. ( 2022 ); Zhang et al. ( 2023c ) ; HA T -P-33 b: Guilluy et al. ( 2023 ) ; HA T -P-49 b: Guilluy et al. ( 2023 ) ; HA T -P-57 b: Hartman et al. ( 2015 ); Stangret et al. ( 2022 ) ; HA T -P-67 b: Bello-Arufe et al. ( 2023 ); Gully-Santiago et al. ( 2024 ) ; HA T -P-70 b: Bello-Arufe et al. ( 2022 ) ; L TT9779 b: Jenkins et al. ( 2020 ); V issaprag ada et al. ( 2024b ) ; MASCARA-2 b: Lund et al. ( 2017 ); T alens et al. ( 2018 ); Casasayas-Barris et al. ( 2019 ) ; T OI-1136 d: Dai et al. ( 2023a ) ; T OI-1235 b: Krishnamurth y et al. ( 2023 ) ; T OI-1268 b: Šubjak et al. ( 2022 ) ; T OI-1431 b: Addison et al. ( 2021 ); Stangret et al. ( 2022 ) ; T OI-1683 b: Zhang et al. ( 2023b ) ; T OI-1728 b: Kanodia et al. ( 2020 ) ; T OI-1807 b: Nardiello et al. ( 2022 ); Gaidos et al. ( 2023 ) ; T OI-2018 b: Dai et al. ( 2023b ) ; T OI-2046 b: Kabáth et al. ( 2022 ) ; T OI-2048 b: Ne wton et al. ( 2022 ) ; T OI-2076 b: Osborn et al. ( 2022 ); Zhang et al. ( 2023b ); Gaidos et al. ( 2023 ) ; T OI-2134 b: Zhang et al. ( 2023a ) ; T OI-2136 b: Ka w auchi et al. ( 2022 ) ; T OI-3757 b: Kanodia et al. ( 2022 ) ; TRAPPIST -1 b: Gillon et al. ( 2017 ); Krishnamurth y et al. ( 2021 ) ; TRAPPIST -1 e: Gillon et al. ( 2017 ); Krishnamurth y et al. ( 2021 ) ; TRAPPIST -1 f: Gillon et al. ( 2017 ); Krishnamurth y et al. ( 2021 ) ; NGTS-5 b: V issaprag ada et al. ( 2022b ) ; K2-25 b: Mann et al. ( 2016 ); Gaidos et al. ( 2020b ); Rockcli ff e et al. ( 2021 ) ; K2-77 b: Gaidos et al. ( 2017 ) ; K2-100 b: Stef ansson et al. ( 2018 ); Barragán et al. ( 2019 ); Gaidos et al. ( 2020a ) ; K2-105 b: Guilluy et al. ( 2023 ) ; K2-136 c: Mann et al. ( 2018 ); Gaidos et al. ( 2021 ); Mayo et al. ( 2023 ) ; KEL T -9 b: Y an & Henning ( 2018 ); Nortmann et al. ( 2018 ) ; K epler -25 c: Guilluy et al. ( 2023 ) ; K epler -68 b: Guilluy et al. ( 2023 ) ; V1298T au c: Da vid et al. ( 2019 ); Suárez Mascareño et al. ( 2021 ); Feinstein et al. ( 2021 ); V issaprag ada et al. ( 2021 ); Schla win et al. ( 2021 ); Gaidos et al. ( 2022 ) ; V1298T au b: Da vid et al. ( 2019 ); Suárez Mascareño et al. ( 2021 ); V issaprag ada et al. ( 2021 ); Gaidos et al. ( 2022 ) ; W ASP-11 b: Allart et al. ( 2023 ) ; W ASP-12 b: Jensen et al. ( 2018 ); Kreidber g & Oklop ˇci 'c ( 2018 ); Czesla et al. ( 2024 ) ; W ASP-39 b: Allart et al. ( 2023 ) ; W ASP-47 d: Guilluy et al. ( 2023 ) ; W ASP-48 b: Bennett et al. ( 2023 ) ; W ASP-52 b: Chen et al. ( 2020 ); Kirk et al. ( 2022 ); V issaprag ada et al. ( 2020 ) ; W ASP-69 b: Nortmann et al. ( 2018 ) ; W ASP-76 b: W est et al. ( 2016 ); Casasayas-Barris et al. ( 2021a ) ; W ASP-77 b: Cortés-Zuleta et al. ( 2020 ), Khalafinejad et al. in prep. ; W ASP-80 b: Salz et al. ( 2015 ); F ossati et al. ( 2022 ) ; W ASP-107 b: Spak e et al. ( 2018 ); Allart et al. ( 2019 ); Kirk et al. ( 2020 ); Spak e et al. ( 2021 ); Hejazi et al. ( 2023 ) ; W ASP-127 b: Lam et al. ( 2017 ); dos Santos et al. ( 2020 ); Allart et al. ( 2020 , 2023 ) ; W ASP-177 b: T urner et al. ( 2019 ); Kirk et al. ( 2022 ); V issaprag ada et al. ( 2022b ) ; W ASP-189 b: Anderson et al. ( 2018 ); Lendl et al. ( 2020 ); Stangret et al. ( 2022 )."}