OPS1
Session assets
Orals: Mon, 19 Sep | Room Andalucia 2
How volatiles were incorporated in the building blocks of planets and small bodies in the protosolar nebula (PSN) remains an outstanding question. Some scenarios consider that planetesimals formed from a mixture of refractory material and volatiles trapped in amorphous ice in the outer nebula, while others hypothesize that volatiles have been incorporated in clathrates or formed pure condensates [1,2]. Here, we study the evolution of volatiles species in the PSN (H2O, CO, CO2, CH4, H2S, N2, NH3, Ar, Kr, Xe and PH3) considering two possible volatiles reservoir in the initial state: amorphous ice (see Figure 1) or pure condensates (see Figure 2). To do so, we use a 1D disk accretion model [3] with radial transport of trace species to compute the radial distribution of volatiles in the PSN. This model includes condensation/sublimation rates of pure condensates, as well as clathration/release rates when enough crystalline water is available. Figure 1 represents the case where volatiles are initially delivered to the PSN in the form of pure condensates. Figure 2 represents the case where volatiles are delivered to the PSN by amorphous ice. Species are released when amorphous grains cross the ACTZ region. Once delivered to the disk, the phase (solid or gaseous) of each species is ruled by the positions of its corresponding condensation and clathration lines. Clathration lines of the considered volatiles are closer to the Sun than their respective condensation lines, except for CO wich have its clathration line further from the sun than its condensation line. Gaseous volatiles condense or become entrapped (depending on the availability of water ice) when diffusing outward the locations of their lines. Conversely, volatiles condensed/entrapped in grains or pebbles are released in gaseous forms when drifting inward their lines. Peaks of abundances form close to each line. Our simulations show that a significant fraction of volatiles can be trapped in clathrates, only if they have initially been delivered in pure condensate forms to the disk. We also show that several regions in the protosolar nebula share a metallicity that is consistent with those measured in the atmospheres of the ice giants [3,4]. These findings have important implications for the formation history of the outer planets
Figure 1 : Scheme showing the disk at initialization and at a given time. The volatiles are initially delivered under pure condensates and vapor. The vapor will condensate into clathrate hydrate, if there is enough crystalline water available. Grains drift inward while vapor undergo diffusion inward and outward. Leading to an accumulation of species at the place of condensation lines.
Figure 2 Scheme showing the disk at initialization and at a given time. The volatiles are initially delivered trapped into amorphous ice and vapor released from amorphous ice at the Amorphous to Crystalline Transition Zone (ACTZ). The clathration line is further than the ACTZ, since there no crystalline water after the ACTZ, clathration cannot happen, or is marginal if the clathration line is close to the ACTZ.
[1] : Gautier, D., Hersant, F., Mousis, O., et al. 2001, ApJL, 550, L227
[2] : Mousis, O., Ronnet, T., & Lunine, J. I. 2019, ApJ, 875, 9.
[3] : Aguichine, A., Mousis, O., Devouard, B., et al. 2020, ApJ, 901, 97.
[4] : Asplund, M., Grevesse, N., Sauval, A. J., et al. 2009, ARA&A, 47, 481.
[5] : Irwin, P. G. J., Toledo, D., Garland, R., et al. 2018, Nature Astronomy, 2, 420.
How to cite: Schneeberger, A., Mousis, O., Aguichine, A., and Lunine, J.: Evolution of the reservoir of volatiles in the protosolar nebula, Europlanet Science Congress 2022, Granada, Spain, 18–23 Sep 2022, EPSC2022-881, https://doi.org/10.5194/epsc2022-881, 2022.
Previous studies of the reflectance spectra of Uranus and Neptune concentrated on individual, narrow wavelength regions, inferring solutions for the vertical structure of gases and aerosols that work well for the wavelength range fitted, but not elsewhere. This has made it difficult to interpret these works with respect to the underlying physical structures. Here we searched for a model of the vertical distribution of haze and methane abundance that fits the whole 0.3 – 2.5 µm range simultaneously, determining more robust solutions.
We reanalysed a set of observations from: HST/STIS (0.3-1.0 µm), IRTF/SpeX (0.8-2.5 µm), and Gemini/NIFS (1.4-1.8 µm). Using a Minnaert limb-darkening model, we were able to partially disentangle cloud opacity and cloud absorbing effects, resulting in a model of the vertical distribution of aerosols and methane that is consistent with the observed reflectivity spectra of both planets and fits all wavelengths simultaneously. This model consists of three main components: Aerosol-1) a deep aerosol layer with a base pressure > 5-7 bar, assumed to be a mixture of H2S ice and photochemical haze; Aerosol-2) a layer of photochemical haze/ice, coincident with a layer of high static stability at the methane condensation level at 1-2 bar; and Aerosol-3) an extended layer of photochemical haze, likely mostly of the same composition as the 1-2-bar layer, extending from this level up through to the stratosphere, the likely origin of the photochemical haze particles. For Neptune, we also need to add a thin layer of micron-sized methane ice particles (Aerosol-4) at ~0.2 bar to explain the enhanced reflection at longer methane-absorbing wavelengths. We suggest that methane condenses onto the haze particles at the base of the 1-2-bar Aerosol-2 layer and forms ice/haze particles that grow very quickly to large size and immediately ‘snow out’, as predicted by Carlson et al. (1988). We suggest that these particles re-evaporate at deeper, warmer levels to release their core haze particles to act as condensation nuclei for H2S ice formation in the Aerosol-1 layer (Fig. 1).
Figure 1. Summary of retrieved aerosol distributions for Uranus (left) and Neptune (right), compared with their assumed temperature/pressure profiles. On each plot are also shown the condensation lines for CH4 (green) and H2S (pink), assuming 10-bar mole fractions of 4% for CH4 and 0.001 for H2S, to move the condensation levels to the approximate levels of the Aerosol-1 and Aerosol-2 layers.
We find that the particles in the 1-2 bar Aerosol-2 layer have the greatest impact on the overall colour of Uranus and Neptune. These particles would appear mostly white in the visible, but become increasingly absorbing at blue/UV and near-IR wavelengths. The Aerosol-2 layer on Uranus has approximately twice the opacity as the same layer on Neptune, giving Uranus a paler blue-green colour than the deeper blue of Neptune, and also explains why Uranus is darker than Neptune at UV wavelengths as can be seen in Figs. 2 and 3.
Figure 2. HST/STIS I/F spectra of Uranus and Neptune averaged along the central meridian of these planets. Also overplotted, for reference are the red, green, blue sensitivities of the human eye.
Figure 3. Modelled appearances of Uranus and Neptune. The right-hand column shows the simulated ‘observed’ appearance of Uranus and Neptune, reconstructed using the fitted Minnaert limb-darkening coefficients extracted from HST/STIS data. The spectra across the discs were convolved with the Commission Internationale de l’Éclairage (CIE)-standard red, green, and blue human cone spectral sensitivities to yield these apparent colours. The preceding columns show how our modelled appearance of these planets changes as we add different components to our radiative-transfer model.
For the deeper Aerosol-1 layer, a darkening of this layer (or possibly, but less likely, a clearing of the aerosols) provides a good match to the observed spectral properties of dark spots on Neptune, such as the Great Dark Spot observed on Neptune in 1989 by Voyager-2 (Fig. 4), and the more recent NDS-2018 spot, first observed by HST/WFC3 in 2018 (Fig. 5). Such spots are seen to have maximum visibility at ~500 nm, but are invisible at UV wavelengths and also at wavelengths longer than 700 nm. We will discuss the dynamical implications these conclusions have on the possible structure of such dark spots.
Figure 4. Representative Voyager-2 ISS images of Neptune, observed in August 1989 in the Clear, UV, Violet, Blue, Green and Orange filters, respectively, of the Narrow Angle Camera (NAC). The Great Dark Spot and Dark Spot 2 are visible, except in the UV channel. Also visible, except in the UV, is the darker belt at 45-55°S.
Figure 5. Observed and reconstructed HST/WFC3 images of Neptune made in 2018. Top: observations with the NDS-2018 feature at upper left and a dark lane at 60°S at lower right. Middle: images reconstructed from our fits to the HST/STIS data, but also including a hole in the deep Aerosol-1 layer (p>5-7 bar) near the central meridian at 15°N, and a clearing at 60°S. Bottom: a second set of images reconstructed from our fits, but where the Aerosol-1 layer is instead darkened near the central meridian at 15°N and at all longitudes at 60°S.
Finally, the homogenous haze layer near the 1-2 bar methane condensation level can be explained by a layer of high static stability caused by the release of latent heat and by the significant decrease with height of the mean atmospheric molecular weight caused by methane condensation, which has a large deep abundance (~4%) and is much heavier than the surrounding H2-He air. Methane condensation would quickly form very large ice particles at the base of this haze layer that immediately ‘snow out’, explaining why a layer of methane ice is not found at its condensation level. Limited mixing from the potential static stability barrier at 1-2 bar could permit external CO entering Neptune’s atmosphere to be trapped in the upper troposphere, removing the need for extreme internal oxygen enrichment in order to explain the wide wings in Neptune’s sub-mm CO lines.
How to cite: Irwin, P., Teanby, N., Fletcher, L., Toledo, D., Orton, G., Wong, M., Roman, M., Pérez-Hoyos, S., James, A., and Dobinson, J.: A holistic aerosol model for Uranus and Neptune, including Dark Spots, Europlanet Science Congress 2022, Granada, Spain, 18–23 Sep 2022, EPSC2022-205, https://doi.org/10.5194/epsc2022-205, 2022.
Uranus and Neptune have a tropopause, stratosphere and a thermosphere inexplicably too hot. This problem is called "Energy Crisis" [1,2]. Radio occultation measurements [3,4] and UV stellar and solar measurements [5,6] by Voyager 2 revealed that their stratosphere reach values of 150 K and their thermosphere reach values of 800 K. Globally-averaged temperature observations [7,8] have also shown a warmer stratosphere than expected. Current radiative equilibrium models can’t explain the observed temperatures [9,10], one or several unknown heat sources are necessarily responsible of this warm atmosphere. Thus, the nature of this or these heat sources is still an open question.
To evaluate the role of several heat sources, we use the radiative-convective model developed by the Laboratoire de Météorologie Dynamique (LMD), previously used for the atmospheres of gas giant planets [11,12]. The model employs the correlated-k formalism to compute gaseous opacities from CH4, C2H2 and C2H6, opacities from collision-induced absorption (H2-H2, H2-He, H2-CH4, He-CH4 and CH4-CH4) assuming a H2 ortho/para fraction at equilibrium. Our model accounts for the internal heat flux, for multiple scattering and Rayleigh scattering. Cloud and aerosols physical properties (exctinction coefficient, single scattering albedo and asymmetry factor) are computed offline assuming Mie scattering, given a particle size distribution and optical constants. In the radiative-convective model, their optical depth is computed for each layer from several entry parameters: the total optical depth (at a reference wavelength), the boundary pressures (bottom and top of the layer), and an aerosol scale height [13].
Three hypotheses were explored separately to induce sufficient heating to warm the atmosphere. We varied the methane mole fraction profile, added an ad hoc aerosol layer, and imposed thermal conduction from the thermosphere.
Values of observed methane molar fraction in the uranian stratosphere being scattered [14], a sensitivity study was carried out by varying the methane homopause altitude and the stratospheric molar fraction. As expected, a high molar fraction makes it possible to heat the stratosphere and to obtain a tropopause altitude similar to observations. Nevertheless, all tests failed to fully heat the tropopause. This is also expected, as much less solar photons in the near infrared reach these levels.
Concerning thermal conduction from the thermosphere, the heating allows to obtain a temperature gradient in upper stratosphere similar to nominal temperature profile from [7]. However, simulated temperature profile is still colder (see figure 1).
Simulations showed that temperature profile of Uranus was very sensitive to the addition of an optically thin aerosols layer confined to a few levels in the stratosphere. To obtain a sufficiently important stratospheric heating, this layer requires that the absorption spectrum is relatively intense in ultraviolet and thermal infrared. This type of spectrum is similar to tholins, black carbon or a carbon-NH3 mixture [15]. The relatively dark spectrum of this layer also induces greenhouse effect in the tropopause.
In the case of Neptune, all hypothesis tested couldn’t warm stratosphere. Therefore, other heat sources are needed to explain observations, such as a dynamical forcing in stratosphere (...) and/or a coupling with thermospheric dynamics and/or auroral heating.
Figure 1: Temperature profiles from our simulations and from observations [3,4,16] at the same latitude and solar longitude of Voyager 2 measurements. Blue curve corresponds to a simulation without aerosols, orange curve aerosolsmodel from [17] is added, and for green curve, we add a conduction from thermosphere with aerosols. Methane molar fraction profile from [14] is used on our simulations.
Acknowledgements
G. Milcareck, S. Guerlet and A. Spiga acknowledge funding from Agence Nationale de la Recherche (ANR) project SOUND, ANR-20-CE49-00009-01.
References
[1] Marley, M. S. and McKay, C. P. (1999). Icarus, 138.
[2] Fletcher, L. N., de Pater, I., Orton, G. S., Hofstadter, M. D., Irwin, P. G. J., Roman, M. T., and Toledo, D. (2020). Space Science Reviews, 216.
[3] Lindal, G. F., Lyons, J. R., Sweetnam, D. N., Eshleman, V. R., Hinson, D. P., and Tyler, G. L. (1987). Journal of Geophysics Research, 92.
[4] Lindal, G. F. (1992). Astronomical Journal, 103.
[5] Herbert, F., Sandel, B. R., Yelle, R. V., Holberg, J. B., Broadfoot, A. L., Shemansky, D. E., Atreya, S. K., and Romani, P. N. (1987). Journal of Geophysical Research, 92.
[6] Broadfoot, A. L., Atreya, S. K., Bertaux, J. L., Blamont, J. E., Dessler, A. J., Donahue, T. M., Forrester, W. T., Hall, D. T., Herbert, F., Holberg, J. B., Hunten, D. M., Krasnopolsky, V. A., Linick, S., Lunine, J. I., Mcconnell, J. C., Moos, H. W., Sandel, B. R., Schneider, N. M., Shemansky, D. E., Smith, G. R., Strobel, D. F., and Yelle, R. V. (1989). Science, 246.
[7] Orton, G. S., Fletcher, L. N., Moses, J. I., Mainzer, A. K., Hines, D., Hammel, H. B., Martin- Torres, F. J., Burgdorf, M., Merlet, C., and Line, M. R. (2014). Icarus, 243.
[8] Fletcher, L. N., de Pater, I., Orton, G. S., Hammel, H. B., Sitko, M. L., and Irwin, P. G. J. (2014). Icarus, 231.
[9] Greathouse, T. K., Richter, M., Lacy, J., Moses, J., Orton, G., Encrenaz, T., Hammel, H. B., and Jaffe, D. (2011). Icarus, 214.
[10] Li, C., Le, T., Zhang, X., and Yung, Y. L. (2018). Journal of Quantitative Spectroscopy and Radiative Transfer, 217.
[11] Guerlet, S., Spiga, A., Sylvestre, M., Indurain, M., Fouchet, T., Leconte, J., Millour, E., Wordsworth, R., Capderou, M., Bézard, B., and Forget, F. (2014). Icarus, 238.
[12] Guerlet, S., Spiga, A., Delattre, H., and Fouchet, T. (2020). Icarus, 351.
[13] Madeleine, J. B., Forget, F., Millour, E., Navarro, T., and Spiga, A. (2012). Geophysical Research Letters, 39.
[14] Lellouch, E., Moreno, R., Orton, G. S., Feuchtgruber, H., Cavalié, T., Moses, J. I., Hartogh, P., Jarchow, C., and Sagawa, H. (2015). Astronomy and Astrophysics, 579.
[15] Rizk, B. and Hunten, D. M. (1990). Icarus, 88.
[16] Lindal, G. F., Lyons, J. R., Sweetnam, D. N., Eshleman, V. R., Hinson, D. P., and Tyler, G. L. (1990). Geophysical Research Letters, 17.
[17] Irwin, P. G. J., Teanby, N. A., Fletcher, L. N., Toledo, D., Orton, G. S., Wong, M. H., Roman, M. T., Perez-Hoyos, S., James, A., and Dobinson, J. (2022). arXiv e-prints.
How to cite: Milcareck, G., Guerlet, S., Montmessin, F., Spiga, A., Leconte, J., Fletcher, L. N., Vatant d'Ollone, J., Roman, M. T., and Lellouch, E.: Radiative heat sources in the stratosphere of Uranus and Neptune, Europlanet Science Congress 2022, Granada, Spain, 18–23 Sep 2022, EPSC2022-1034, https://doi.org/10.5194/epsc2022-1034, 2022.
Neptune's incomplete ring arcs have been stable since their discovery in 1984 by stellar occultation. Although these structures should be destroyed within a few months through differential Keplerian motion, imaging data over the past couple of decades have shown that these structures remain stable.
We present the first SPHERE near-infrared observations of Neptune's ring arcs taken at 2.2 μm (broadband Ks) with the IRDIS camera at the Very Large Telescope (VLT) in August 2016.
We derive accurate mean motion values for the arcs and the nearby satellite Galatea. The trailing arcs Fraternité and Égalité have been stable since they were last observed at VLT in 2007 (cf. Fig 1 left). Furthermore, we confirm the fading away of the leading arcs Courage and Liberté. Finally, we confirm the mismatch between the arcs' position and the 42:43 inclined and eccentric corotation resonances with Galatea, thus demonstrating that no 42:43 corotation model works to explain the azimuthal confinement of the arcs' material (cf. Fig. 1 right).
Fig. 1 (left) 68 projected and co-added images of Neptune’s equatorial plane, revealing material along the arcs Fraternité and Égalité, as well as the satellites Proteus (P) and Galatea (G). The frame is 7.1x10.7 arcsec2 wide. (right) Equivalent width of the arcs (Fraternité and Égalité), at an angular resolution of 2°. The X-axis origin is the longitude LFr of the centre of the Fraternité arc measured from the J2000.0 ascending node, where LFr = 217.43 deg, at the reference epoch. (Souami et al. 2022)
We also present here our NOC21 (Neptune stellar Occultation Campaign 2021) observational campaign to study the evolution of the Neptunian system since Voyager-2.
On October 7th, 2021, we led and organised the largest occultation and imaging campaign by the Neptunian system ever undertaken since the Voyager flyby in 1989. This occultation event was observable across the Americas (North and South) as well as in Hawaii, and the campaign involved all large telescopes in the region that were equipped with fast cameras either in the infra-red K band (2.2 μm) or in the visible using a CH4 filter at 890 nm. Combined with Adaptive Optics images, these observations will bring new constrains on Neptune's rings dynamics
We will use the occultation data collected during the NOC21 campaign along with AO data obtained at Keck observatory around the occultation time, to address the following points:
- the study of Neptune’s arcs’ system evolution using adaptive optics data,
- the detection of the main Le Verrier and Adams rings (see Fig.1) in the NOC21 data to measure their optical depths. In particular, we will search for possible changes in density of the Adams rings since the Voyager-2 mission in 1989,
- the search (in the NOC21 data) for any other narrow rings and/or small satellites and/or local debris such as arcs.
How to cite: Souami, D., Sicardy, B., Renner, S., and Langlois, M.: The evolution of Neptune’s arcs since Voyager-2. VLT/SPHERE observations of Neptune’s ring arcs and the 2021 Neptune stellar occultation campaign, Europlanet Science Congress 2022, Granada, Spain, 18–23 Sep 2022, EPSC2022-1263, https://doi.org/10.5194/epsc2022-1263, 2022.
The five major Uranian satellites provide important constraints on compositional and dynamical evolution in the Uranian system. These bodies experienced endogenic activity in the past, morphologically distinct from those on Saturnian icy moons. Uranian moons also exhibit compositional differences with their Saturnian counterparts. Because our understanding of these bodies is based on the single 1986 Voyager flyby and telescope observations, much remains to be discovered and understood about how these moons formed and evolved.
- 1. Introduction
The five major Uranian moons (diameters 930 to 1575 km and best density estimates of ~1.2 to 1.7 kg/m3)form a regular family of ice-rich planetary bodies, and record accretional conditions in the Uranian system. Voyager 2 mapping of <40% of their surfaces provides our only resolved record of these worlds [1,2]. Earth- and space-based telescope observations provide additional constraints. Here we review these moons in context with the Saturnian system, with a view towards the now likely return to the Uranian system as a major NASA Flagship.
Figure 1: Voyager views of the 5 major Uranian icy moons.
- 2. Composition and Interiors
Uranian satellites are markers of the composition of the region they formed in. Telescope observations indicate surface compositional differences from Saturnian moons. Saturnian moons have bright surfaces with strong H2O ice absorption bands (except Iapetus’ dark), partly due to E-ring deposition [e.g., 3]. Uranian moons have darker surfaces of H2O-ice mixed with spectrally neutral components possibly rich in organics and/or hydrated silicates. CO2 on Saturnian moons is complexed with other species (except CO2-ice on Enceladus) [4,5]; CO2 is in crystalline form on Uranian moons [4,5].
Limited topography indicates Uranian moons are more rugged than inner Saturnian moons [7], provisionally attributed to global thermal resetting in saturnian moons. Enceladus is the only confirmed ocean world among these moons (deep oceans have been suggested at Dione [8] and Mimas [9]). Saturnian moons are involved in mean motion resonances that drive obliquity or eccentricity, whereas Uranian moons are not but likely shared such resonances in the past [10,11].
Radioactive heating estimates range from ~0.1 mW/m2 for Miranda (similar to Mimas) to ~0.7 mW/m2 for Ariel/Umbriel (similar to Dione/Rhea, but three orders magnitude less than measured at Enceladus) to ~1.1 mW/m2 for Titania/Oberon [12]. If Uranian moons preserve liquid, it is likely as residual oceans < 30 km thick, or brine-filled porous cores [12]. Miranda is unlikely to preserve liquid presently unless recently tidally heated.
- 3. Tectonics, Volcanism, Past Heat Flow
The five Uranian moons betray various levels of tectonic activity, from deep walled troughs of Ariel (Fig. 2) to cryptic deep features on Umbriel’s limb [2,7]. Some are ancient but global extent and associated stresses are uncertain due to mapping limits. (Cryo)volcanism of higher relief than in the Saturnian system is widespread on Miranda and Ariel (and unresolved on the others) but composition and eruption state are poorly known [7].
Figure 2. Voyager stereo-derived DEM showing tectonic and resurfacing features on Ariel. Cylindrical map projection.
Ariel’s estimated past heat flux (6–92 mW/m2) [13,14] overlaps Tethys, Dione, Rhea, Triton, Pluto, and Charon (Fig. 3), and are consistent with tidal heating from past resonances [13,14]. Miranda’s maximum estimated heat flux (31–140 mW/m2) [14,15] is higher than Ariel, Tethys, Dione, and Rhea, and overlaps estimates for Enceladus [16,17]. High heat fluxes for Miranda suggest increases to eccentricity from past resonances [10,11].
Figure 3. Comparison of heat flux estimates on icy bodies.
- 4. Comparison and Exploration
Like Saturn’s icy moons, the five major Uranian satellites are diverse, from little geologic activity to near global resurfacing as little as ~0.1 Gyr ago [17]. The two satellite systems are not identical, representing formation in distant compositionally distinct part of the Solar System, and under possibly much different formation conditions. Cassini showed the value of a dedicated orbiter with instrumentation across the spectrum. Aside from resurfacing of Dione, Tethys, and esp. Enceladus [18], Cassini revealed magnetospheric and ring particle surface alterations [18, 19], satellite ring deposition [18,20] and other processes. Whether these occur(red) in the Uranian system must await a dedicated Uranian orbiter. Such observations and interior property constraints from close encounters, will allow better understanding of formation processes of these moons and the Uranian system.
References
[1]. Smith, B., and 30 others, Science, 233, 43-64, 1986.
[2]. Croft, S., and L. Soderblom. In Uranus, Univ. Arizona Press, 1990.
[3]. Verbiscer, A., R. Frech, M. Showalter, and P. Helfenstein, Science, 315, 815, 2007.
[4]. Cartwright, R.J., Emery, J.P., Rivkin, A.S., Trilling, D.E. and Pinilla-Alonso, N., Icarus, 257, pp.428-456, 2015.
[5]. Cartwright, R.J., Emery, J.P., Pinilla-Alonso, N., Lucas, M.P., Rivkin, A.S. and Trilling, D.E., Icarus, 314, pp.210-231, 2018.
[6] Schenk, P, White O, Byrne P, Moore J. 2018. In Enceladus and the icy moons of Saturn (ed. P. Schenk), pp. 237–265. Tucson, AZ: University Arizona Press, 2018.
[7] Schenk, PM, and J.M. Moore, Phil. Trans. R. Soc. A. 378, 20200102, 2020
[8] Beuthe, M., et al., Geophys. Res. Lett. 43, 10, 088–10, 2016
[9] Rhoden, A., and M. Walker, Icarus, 376, 114872, 2022
[10] Tittemore, W. and Wisdom J., Icarus 85, 394–443. 1990
[11] Cuk, M., El Moutamid, M., & Tiscareno, M. S. Planetary Science Journal, 1, 22, doi: 10.3847/PSJ/ab9748, 2020
[12] Castillo-Rogez, J. C., and 7 others, submitted to J. Geophys. Res., 2022; AGU Fall meeting, 2021.
[13] Peterson, G, Nimmo F, Schenk P., Icarus 250, 116–122, 2015
[14] Beddingfield, C., Cartwright R., Leonard E., Nordheim T., Scipioni F., 2022 PSJ 3, 106, 2022
[15] Beddingfield, CB, Burr DM, Emery JP., Icarus 247, 35–52, 2015
[16] Bland M. T., Beyer R. A. and Showman A. P., Icarus, 192, 92, 2007; Bland M. T., Singer K. N., McKinnon W. B. and Schenk P. M. 2012 Geophys. Res. Lett., 39, L17204, 2012
[17]. Kirchoff, M., et al, Planet Sci. J., 3, 42, 2022.
[18]. Schenk, P., D. Hamilton, R. Johnson, W. McKinnon, J. Schmidt, M Showalter, Icarus, 211, 740-757, 2011.
[19] Howett, C. et al., Icarus, 216, 221-226, 2011.
[20] Ip, W., Geophys. Res. Lett., 33, L16203, 200.
How to cite: Schenk, P., Cartwright, R., Beddingfield, C., and Castillo-Rogez, J.: Uranian Mid-Sized Icy Moons: Unique Targets of Exploration in the Outer Solar System, Europlanet Science Congress 2022, Granada, Spain, 18–23 Sep 2022, EPSC2022-1245, https://doi.org/10.5194/epsc2022-1245, 2022.
Posters: Mon, 19 Sep, 18:45–20:15 | Poster area Level 1
Background: Some fundamental properties of the interiors of the Ice giants Uranus and Neptune are far from being understood. According to structure models which follow observed gravitational harmonics J2 and J4, the interiors are composed of a H/He-rich envelope transitioning into an ice-rock interior. Formation theories can explain their current structure given certain conditions related to the formation location within the protoplanetary disk nebula and the degree of gas depletion during said formation are met, conditions which themselves seek an explanation [1].
A rather sharp boundary yields the simplest solution to the gravity field [2] and the luminosity of the planets [3]. Recently, the phase separation of two major constituents (water and molecular hydrogen) in the evolved and cooled planets Uranus and Neptune has been proposed as an explanation for the presence of this sharp interface. Furthermore, different H2O/H2 demixing states may offer an explanation for the paradox between intrinsic heat flux of both planets [2]. On the other hand, evolution models guided by the observed luminosities suggest the existence of a compositional barrier inhibiting or slowing down convection. Assuming small differences in the structure of both planets, this mechanism can account for the faintness of Uranus and the brightness of Neptune [3]. However, their evolution could be affected by possible complementary processes such as phase separation [2] or condensation in their atmospheres.
In this work, we follow up on the possibility of demixing between the major constituents in the H-He-H2O system. First, we show that interior models which adjust to the luminosity [3] lead to temperatures above the critical temperature for H2O/H2 demixing to occur, and thus would predict a protosolar atmospheric helium abundance. Second, we adopt the assumption of an initially more homogeneous interior which did cool sufficiently to allow for H2O/H2 demixing [2]. We find that for deep interior H/He phase separation, which occurs at higher temperatures, favourable interior conditions were met much earlier in the evolution [4].
Figure 1: Helium particle fraction as a function of assumed water mass fraction in the deep interior of Uranus. Error bars placed at ZH2O= 0 show the observed atmospheric He abundances of Jupiter, Saturn, Uranus, and Neptune, while also a value for Jupiter's deep interior. The grey area delimits the uncertainty region marking the lower limit of He/H ratio necessary for phase separation to occur in Uranus. This figure shows how the He/H ratio relevant to H/He demixing behaves with the water abundance.
In Figure 1, we plot the He particle fraction as a function of the assumed water mass fraction in the deep interior assuming full dissociation of water and hydrogen. Protosolar values for the atmospheric helium abundance Y were adopted for Uranus (0.26-0.28). The grey region shows the uncertainty region above which H/He demixing can occur, inferred from the H/He non-ideal entropy phase diagram of [4]. Figure 1 shows H/He phase separation cannot take place in a highly ice-rich interior ZH2O greater ~0.85. However, moving towards lower water contents, demixing can occur at pressure levels between 1.5 to 4 Mbar [4]. The atmospheric helium abundance prediction we aim for could be compared to probed in-situ measurements in the future.
Acknowledgements:
The authors acknowledge support from the Research Unit 2440/2 funded by the DFG (Deutsche Forschungsgemeinschaft).
[1] Frelikh R. and Murray-Clay R.A., AJ 154:98 (2017). [2] Bailey E. and Stevenson D.S., PSJ 2:64 (2021). [3] Scheibe L. et al., A&A 650:A200 (2021). [4] Schöttler M. and Redmer R., PRL 120:115703 (2018).
How to cite: Cano, M., Nettelmann, N., and Tosi, N.: On the possibility of phase separation in the H-He-H2O system in the Ice Giants, Europlanet Science Congress 2022, Granada, Spain, 18–23 Sep 2022, EPSC2022-1209, https://doi.org/10.5194/epsc2022-1209, 2022.
Water vapour in the stratospheres of Uranus and Neptune has previously been shown to originate from external sources. These sources could include comet impacts [4], interplanetary dust particles [8], or rings and moons [1]. Stratospheric water was first detected on Uranus and Neptune by the Short-Wavelength Spectrometer (SWS) on the Infrared Space Observatory (ISO) [2], but the uncertainties were relatively large due to lack of constraint on the vertical water profiles and relatively low spectral resolution of the observations.
Here we present new observational constraints on Uranus’ and Neptune’s externally sourced stratospheric water abundance using disc-averaged high spectral resolution observations of the 557 GHz water emission line from Herschel’s Heterodyne Instrument for the Far-Infrared (HIFI). On both planets the emission line is significantly broadened by disc-averaging of Doppler shifts from planetary rotation, which was carefully accounted for in our analysis [10]. Derived stratospheric column water abundances are 0.56+0.26-0.06 x 1014 cm-2 for Uranus and 1.9+0.2-0.3 x 1014 cm-2 for Neptune. These results imply Neptune has about four times as much stratospheric water as Uranus, and are consistent with previous determinations from ISO-SWS and Herschel-PACS, but with improved precision.
For Uranus excellent observational fits are obtained by scaling photochemical model profiles [3,7] or with step-type profiles with water vapor limited to <=0.6mbar. However, Uranus’ cold stratospheric temperatures imply a ~0.03mbar condensation level, which further limits water vapor to pressures <=0.03 mbar. Neptune’s warmer stratosphere has a deeper ~1 mbar condensation level, so emission-line pressure broadening can be used to further constrain the water profile. For Neptune, excellent fits are obtained using step-type profiles with cutoffs of ~0.3-0.6 mbar or by scaling a photochemical model profile [7]. Step-type profiles with cutoffs >=1.0 mbar or <=0.1 mbar can be rejected with 4σ significance. Rescaling photochemical model profiles from [7] to match our observed column abundances implies similar external water fluxes for both planets: 8.3+4.0-0.9 x 104 cm-2s-1 for Uranus and 12.7+1.3-2.0 x 104 cm-2s-1 for Neptune.
This inferred water influx rates suggest that Uranus and Neptune may in fact have very similar IDP fluxes, unless there are significant water-loss processes that are not accounted for in current photochemical models [3,7]. This is unexpected as the IDP flux on Neptune is expected to be higher due to its closer proximity to the Kuiper belt. For example, the dynamical model of [8] predicts that the flux of IDP grains is around seven times higher on Neptune than on Uranus, but model uncertainties are large enough so as not to preclude a similar flux. The comet impact rates on Uranus and Neptune are also predicted to be quite similar [5,11], so both planets may experience similar external flux processes.
Our new analysis suggests that Neptune’s approximately four times greater observed water column abundance is primarily caused by its warmer stratosphere preventing loss by condensation, rather than by a significantly more intense external source. Larger error bars on the Uranus estimates are due to greater uncertainty in the high-altitude temperature profile. To reconcile these water fluxes with other observed stratospheric oxygen species (CO and CO2) requires either a significant CO component in interplanetary dust particles (Uranus) or contributions from cometary impacts (Uranus, Neptune). In particular, the large CO abundance in Neptune’s stratosphere suggests that we just happen to be observing Neptune at a time shortly after a large comet impact [4,6,9].
Further details of our results and analysis are available in our recent publication [10].
Fig1: Herschel-HIFI line-to-continuum ratio spectra of the 557GHz water line for HRS and WBS spectrometers. The water line is clearly visible at high signal to noise on both planets, but the line is broadened due to Doppler shift combined with the disc-broadened nature of the HIFI spectra. (Figure from https://doi.org/10.3847/PSJ/ac650f, see reference [10]).