Exoplanet atmospheres

Abstract:

Saturn’s largest moon Titan has a nitrogen-methane atmosphere and a rich organic photochemistry. Dissociation of Titan’s molecular methane and nitrogen into N and methyl (CH3) radicals forms the basis of this photochemistry and results in a vast array of hydrocarbon and nitrile species. The abundance of CH3 is thus of critical importance to understanding Titan’s atmospheric chemistry. CH3 is predicted by photochemical models and must be present to explain Titan’s trace gas composition, but has never been directly observed. Cassini’s mass spectrometer was unable to make a detection as the extreme reactivity of radicals results in reactions on the instrument wall (e.g. recombination with H) before detection is possible. Emission features in the infra-red are also very weak, so detection from remote-sensing spectroscopy has previously not been possible. Here we use the very high sensitivity of the James Webb Space Telescope’s (JWST) Mid-InfraRed Instrument (MIRI) to detect emission from CH3 at 16.5 microns. We have used this to validate model predictions that underpin Titan’s rich atmospheric chemistry.JWST/MIRI observations were taken in Medium Resolution Spectroscopy (MRS) mode on 11th July 2023 as part of Guaranteed Time Observation programme 1251 [Nixon et al., 2025]. Observations were reduced using the standard pipeline and combined to give a disc-averaged spectrum (Fig 1). The observed spectrum was compared to a forward model generated with a reference Titan atmosphere using the NEMESIS radiative transfer suite [Irwin et al., 2008]. The reference atmospheric temperature profile was based on observation from Cassini half a Titan year previous, augmented with ground-based measurements from ALMA and in-situ measurements from the Huygens probe (Fig 2a). A baseline atmospheric composition was compiled from Cassini/Huygens measurements [Teanby et al., 2019]. For the CH3 profile, in the absence of measurements, we used the predicted abundance from a photochemical model [Vuitton et al., 2019] (Fig 2a).The abundance profile of CH3 is expected to be extremely steep with very high fractional abundances in the thermosphere (100 ppm at 1000km) and much lower abundances in the stratosphere and mesosphere (1 ppb at 300km). Peak emission under conditions of local thermodynamic equilibrium should originate from the mid-thermosphere at an altitude of ~800km (Fig 2b). However, our analysis shows that non-local thermodynamic equilibrium (non-LTE) emission is expected due to very low thermospheric pressures [Nixon et al., 2025]. This supresses emission below that expected from the Planck function and reduces infra-red emission from thermospheric CH3 to negligible levels. When non-LTE effects are considered, we find that the emission instead originates from the stratopause region (~300km) where CH3 abundances are predicted to be around 1 ppb (Fig 2c).Agreement between forward modelled non-LTE emission using the photochemical model profile and the JWST/MIRI observation match very well (Fig 1) – confirming the model predicted abundances are consistent with conditions in Titan’s middle atmosphere. Our initial results were presented in Nixon et al., (2025). Here we present an updated analysis using improved pipeline processing, more in-depth treatment of the disc-averaged nature of the observation, and provide formal limits on the CH3 abundance profiles. The consistency of our results with predictions from photochemical models gives confidence to current chemical schemes for Titan’s low-order chemistry, which provides a sound basis for a deeper analysis of Titan’s more exotic species such as high-order hydrocarbons and poly-aromatic hydrocarbons.ReferencesIrwin, P.G.J., et al., 2008. The NEMESIS planetary atmosphere radiative transfer and retrieval tool. Journal of Quantitative Spectroscopy and Radiative Transfer 109, 1136–1150.Nixon, C.A., et al., 2025., Titan’s Atmosphere in Late Northern Summer from JWST and Keck Observations. Nature Astronomy, in press.Teanby, N.A., et al., 2019. Seasonal Evolution of Titan’s Stratosphere During the Cassini Mission. Geophysical Research Letters 46, 3079–3089.Vuitton, V., et al., 2019. Simulating the density of organic species in the atmosphere of Titan with a coupled ion-neutral photochemical model. Icarus 324, 120–197.Fig 1: JWST/MIRI disc-average spectrum compared with forward models with and without CH3. The model including CH3 provides a much better fit to the observations.Fig 2: (a) Titan’s atmospheric temperature structure and uncertainty envelope from Nixon et al. (2025), along with photochemical model prediction of the CH3 profile from Vuitton et al. (2019). (b) Contribution functions for LTE case with nominal temperature profile (green), hot temperature limit (red) and cold temperature limit (blue). For LTE, peak emission would be from the thermosphere at ~800km, but this is not realistic. (c) Contribution functions for a more realistic non-LTE emission case peak at ~300km around the mesopause as non-LTE effects suppress emission at very low pressures. Our observations are thus most sensitive to abundances around the stratopause. 

Abstract:

Radiative transfer analyses of spectra obtained from Uranus and Neptune have revealed the presence ofa cloud layer at pressures greater than ~2 bar (1,2). The detection of hydrogen sulfide (H₂S) gas abovethis cloud layer on both planets (3,4) suggests that H₂S ice is the most likely main constituent. Thisinterpretation is further supported by the expectation that methane (CH₄) clouds condense at higheraltitudes (5). However, due to their depth and observational limitations, our understanding of theproperties of H₂S clouds on these planets remains very limited.To investigate the properties of H₂S clouds in the atmospheres of Uranus and Neptune, we employed aone-dimensional cloud microphysics model originally developed for Titan and Mars (6,7). The modelincludes nucleation, condensation, evaporation, coagulation, and precipitation processes, and haspreviously been used to simulate haze and CH₄ cloud microphysics in the Ice Giants (5,8,9).Figure 1 shows, as an example, simulated H₂S ice profiles for Uranus using this microphysical model.The vertical transport of H₂S gas is simulated using an eddy diffusion coefficient (Keddʏ), which controlsthe supply of vapor for cloud nucleation and particle growth. We employed the Keddʏ profiles derivedin [10] for H₂S abundances of 10× and 30× solar. Since several cloud microphysical parameters for H₂Sremain uncertain (e.g., the contact parameter), different values are tested in the simulations. In theexample shown, the model indicates cloud bases near 5.3 bar for 10× solar abundance and 6.4 bar for30× solar. Near the cloud base, particle mean radii range from 40 to 55 μm, depending on the assumedcontact parameter and abundance. At higher altitudes, particle sizes decrease; for instance, at ~3 bar,mean radii are around 20 μm. In general, H₂S cloud simulations produce higher opacities than CH₄clouds.In this work, we will present a series of cloud microphysical simulations of H₂S clouds in the Ice Giants.Various cloud properties, such as particle size distributions and precipitation rates, will be constrained.We will also discuss the implications of our results for the atmospheric circulation of these planets andfor the future exploration of Uranus.Figure 1. Vertical distributions of H2S ice (g/m³) for Uranus, simulated for different values of the cloudcontact parameter and deep H2S abundances. These simulations employ the Keddʏ profiles calculated in[10] for the corresponding H2S abundances.References: [1] P. G. Irwin, et al., JGR: Planets, 127, e2022JE007189. [2] L. Sromovsky, et al., Icarus,Volume 317, (2019) [3] P. G. Irwin, et al., Nature Astronomy 2, 420 (2018). [4] P. G. Irwin, et al.,Icarus 321, 550 (2019). [5] D. Toledo, et al., A&A, 694, A81 (2025). [6] P. Rannou, et al., Science 311,201 (2006). [7] F. Montmessin, et al., JGR: Planets 107, 4 (2002). [8] D. Toledo, et al., Icarus, 333, 1-11, (2019). [9] D. Toledo, et al., Icarus, Volume 350, (2020). [10] H. Ge, et al., The Planetary ScienceJournal,5, 101(2024). 

Abstract:

In 2018, analysis of Gemini-NIFS near-infrared observations revealed the probable presence of H2S above the main cloud deck on Neptune [1]. The spectral signature of the gas was found to be much stronger at Neptune's south pole compared to regions nearer the equator.Conversely, analysis of Neptune's microwave emission with ALMA suggested strongly enhanced H2S abundances at midlatitudes [2], with much less at the south pole. Determining the true variation of H2S with latitude is crucial for understanding the tropospheric circulation of Neptune.We present our analysis of observations of Neptune from VLT-SINFONI in 2018. Using a limb-darkening approximation, we are able to fit the reflected solar radiance from multiple zenith angles, which allows us to discriminate between gas and aerosol opacity. Despite the lower spectral resolution of this instrument compared to Gemini-NIFS, we are able to detect the H2S spectral signature. With our radiative transfer retrieval code, archNEMESIS [3], we use nested sampling to fit a parameterised cloud model (similar to that of [4]) to these observations over a range of latitudes. We prescribe a latitudinally varying deep methane abundance derived from recent VLT-MUSE observations [5], which enables us to constrain the depth of the cloud top.Our retrieved results are in agreement with the results derived from ALMA [2] - we find a significant enhancement of deep H2S at Neptune's southern midlatitudes, decreasing towards the equator and the pole. Our results show a much deeper cloud top towards the pole, resulting in the increased cloud top column abundance of H2S observed here in the previous near-infrared analysis [1].Figure 1: A comparison of fits to a spectrum extracted from the 50°S to 60°S latitude band, with a model including H2S (blue) and a model without H2S (red). Note the significant discrepancy around 1.58 microns. The models are fitted to spectra at two zenith angles simultaneously.[1] Irwin, P. G., Toledo, D., Garland, R., Teanby, N. A., Fletcher, L. N., Orton, G. S., & Bézard, B. (2019). Probable detection of hydrogen sulphide (H2S) in Neptune’s atmosphere. Icarus, 321, 550-563.[2] Tollefson, J., de Pater, I., Luszcz-Cook, S., & DeBoer, D. (2019). Neptune's latitudinal variations as viewed with ALMA. The Astronomical Journal, 157(6), 251.[3] Alday, J., Penn, J., Irwin, P. G., Mason, J. P., & Yang, J. (2025). archNEMESIS: an open-source Python package for analysis of planetary atmospheric spectra. arXiv preprint arXiv:2501.16452.[4] Irwin, P. G., Teanby, N. A., Fletcher, L. N., Toledo, D., Orton, G. S., Wong, M. H., ... & Dobinson, J. (2022). Hazy blue worlds: a holistic aerosol model for Uranus and Neptune, including dark spots. Journal of Geophysical Research: Planets, 127(6), e2022JE007189.[5] Irwin, P. G., Dobinson, J., James, A., Wong, M. H., Fletcher, L. N., Roman, M. T., ... & de Pater, I. (2023). Latitudinal variations in methane abundance, aerosol opacity and aerosol scattering efficiency in Neptune's atmosphere determined from VLT/MUSE. Journal of Geophysical Research: Planets, 128(11), e2023JE007980.

Abstract:

IntroductionJupiter’s northern Equatorial Zone (EZn) and southern North Equatorial Belt (NEBs) are dominated by three features: five-micron hotspots (seen as North Equatorial Dark Features, NEDFs, in the optical), white cloud plumes, and complex local circulation. These features are influenced by the NEBs jet, which is modulated by a meridionally trapped Rossby wave, in conjunction with the high concentration of ammonia in the EZ and the ammonia depletion in the NEB. Numerous measurements have been made of the temperature, aerosol, and ammonia distributions in this region (c.f. Fletcher et al., 2020). And a number of models have been partially successful at explaining the interrelationships between the observed features (c.f Showman & Dowling, 2000). Here we explore the ammonia and cloud height distribution during 2024-25, when NEDFs and five-micron hotspots were prominent, using the optical band-average technique (Hill et al., 2024, Irwin et al., 2025). We show that while many sensing methods highlight the ammonia and aerosol depletion in five-micron hotspots, this band average method highlights enhancements in ammonia to the south of the hotspots.ObservationsMultiple observations on 2025-01-06 were made allowing coverage of a wide range of longitudes and coverage of a given longitude at several zenith angles. Figure 1 shows maps constructed using the method of Hill et al. (2024). An empirical limb correction is applied in addition to a weighted averaging scheme for overlapping observations. The data clearly show that enhanced ammonia regions lie to the south of NEDFs (labeled 1-4 in order of ascending longitude). For the ammonia enhancements we observe a planetary wave number of nine, within the range of hotspot and NEDF wavenumbers typically observed.DiscussionThe NEBs jet speed peaks at about  7° N, which in fact marks the boundary between the NEDFs and the ammonia enhancements. Anticyclonic gyres are a known feature seen in the same location as we show ammonia enhancements (c.f. Choi et al., 2013). We hypothesize that these gyres are regions of uplift and outflow, bringing up ammonia rich air from deeper levels of the atmosphere. The NEDFs are thought to be areas of subsidence, with cyclonic flow, where dryer air descends from above and results in a clearing of aerosols. Figure 1D shows this schematically with upwelling occurring at the gyres, horizontal winds carrying condensates from the upwelling source to the east and northeast as the visible cloud plumes, and descending clear air in the NEDFs.To further support this hypothesis, we analyze the ammonia mole fraction and cloud pressure at the NEDFs, gyres, and in the plumes through a regions-of-interest (ROI) approach. Figure 2 shows a longitudinal subset of the data in Figure 1, focusing on ammonia regions 3 and 4. Rectangles outline the ROIs which are analyzed for three observation times in Figure 2A. Figure 2B shows a time series of average values at each observation time for cloud pressure and ammonia mole fraction along with statistical errors. Finally, 2C shows scatter plots of the average cloud pressure versus the ammonia abundance. Note the very clear clustering of points where the NEB sample provides a consistent reference with relatively high pressure and very low ammonia abundance. Following the upwelling ammonia, eastward advection of plume aerosols, and NEDF subsidence from Figure 1, we can trace an ammonia cycle between its gaseous source and sink, with an intermediary aerosol state.Future WorkHundreds of observations of NEDFs and ammonia enhancements in the EZn have been made in 2024-25 using the Hill et al. (2024) technique. This data set will be analyzed and assessed for the statistical consistency of the results presented here. In addition, this data set will be compared to complementary multispectral observations to help discriminate why the optical method seems to so clearly detect ammonia enhancements at the 1-2 bar pressure level and why these enhancements appear broad enough to overlap NEDFs.Figure 1. Ammonia mole fraction, cloud pressure, and visual context maps created from observations on 2025-01-06 using an 11 inch Schmidt-Cassegrain telescope. A) Ammonia mole fraction (ppm) with enhanced areas labeled 1-4 in order of ascending longitude. The black circle at left shows the approximate spatial resolution of the data. B) Cloud pressure (mbar). C) Visual context image with selected contour overlays to show enhanced ammonia mole fraction and lowest pressure (highest) clouds. D) Same as C), but with arrows indicating presumed upwelling (black ⊙), downwelling (white ⦻), and horizontal flow (red arrows). Figure 2. Two ammonia enhancements (4 & 3 from Figure 1), associated plumes, and NEDFs are analyzed for cloud pressure and ammonia abundance. Three observations are assessed with the targets near nadir viewing. A) Ammonia mole fraction, cloud pressure, and visual context image with overlaid rectangles indicating regions-of-interest (ROIs). B) Time series of cloud pressure (left) and ammonia mole fraction (right) over the three observations. C) Scatter plot of all ammonia and cloud measurements in each ROI (left) and of the averages over the three observations. Note that the NEB data are provided as a stable reference.ReferencesChoi, D. S. et al. 2013. Icarus, 223, 832. Hill, S. M. et al. 2024. Earth and Space Science, 11(8), e2024EA003562.Fletcher, L. N. et al. 2020. Journal of Geophysical Research (Planets), 125, e06399. Irwin, P. G. J. et al. 2025. Journal of Geophysical Research: Planets, 130(1), e2024JE008622. Showman, A. P., & Dowling, T. E. 2000. Science, 289, 1737-1740. 

Abstract:

Survive or not survive, that is the question of the 500-hour JWST Rocky Worlds DDT Program. Whether a terrestrial planets’ atmosphere can suffer under the intense XUV of its host, or if it completely escapes, these are the questions we explore. Zahnle & Catling (2017) defined the Cosmic Shoreline, but recent observations from JWST reveal airless worlds around M-stars, calling for a refinement of this “receding” shoreline (Pass et al. 2025). M-stars spend a longer time in pre-main sequence, subjecting their orbiting worlds to some higher intensity XUV activity. This complicates our present understanding of this shoreline. Investigating chemical effects of planet-star interactions could be the key to a more complete picture of this shoreline.  We investigate the interplay between photochemistry, mixing, and escape of carbon dioxide atmospheres under intense and mild XUV fluxes as follow on work to both Johnstone et al. (2018) and Nakayama et al. (2022). We expand on this work by adopting thermal structure models from Nakayama et al. (2022) and apply them to identify key chemical pathways for escape. We create a reduced C-O chemical network including neutral and ionic species to identify these pathways. As photochemistry simulations take into account many reactions, these 1D calculations are too computationally expensive to be done in 3D. Although rudimentary at best, the mixing parameter– eddy diffusion term, K_zz, comprises the dynamical element of 1D photochemical simulations. Here, we consider the mixing of photochemical products in competition with escape to explore the chemical pathways of retention and loss. We compare the photochemical model results for active and inactive cases for the Trappist-1 system planets. Then, using the resulting composition-dependent heating and cooling rates for Trappist-1 planets, we assess their propensity for efficient atomic line cooling versus escape. We follow the work of Chatterjee & Pierrehumbert (2024) in this assessment.  Finally, using our pathway analysis, we find an analytical formula for calculating an energy-limited escape boundary for these planets based on composition.  It is important here to note the limitations of 1D work. First, there exists an exchange of rigor between modelling chemistry and dynamics. Insights from this work are ripe for implementation into 3D GCMs, especially in response to incorporating UV-driven processes for thermospheric modelling mentioned in Ding and Wordsworth (2019). Second, interaction with the interior is important in the early phase of planetary formation, i.e., the magma ocean phase. Due to exchange between atmosphere and magma early in the planet’s formation, incorporation with an interior-atmosphere model would better constrain higher pressure chemical abundances. Although this work focuses on the upper atmosphere, extrapolation to the surface environment is a key goal for understanding a planet.  Considering planet-star interaction is imperative for the selection of targets for observation. However, it is also important when considering anomalous detections of atmospheres around planets predicted to not have an atmosphere. This could be a first step in determining an atmosphere as non-primary and/or distinguishing between an airless planet and one with high altitude haze.