adsabs/AstroCLIMB huggingface dataset:
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Variability of overall mass influxes of meteoroid impacts, in g cm−2 s−1, as a function of the selenographic longitude and latitude determined by the MIF for four Moon phases: new moon (Day 0; panel (a)), first quarter (Day 8; panel (b)), full moon (Day 15; panel (c)), and third quarter (Day 22; panel (d)). Note that the corresponding fluxes are plotted over an elevation map of the Moon centered on the prime meridian. The vertical black lines represent the positions of the antihelion (LT = 0 hr, nightside), apex (LT = 6 hr, dawn terminator), and helion (LT = 12 hr, noon side) directions.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Panel (a) represents the total melting rates, in g cm−2 s−1, as a function of the selenographic latitude and over a lunar month. Panel (b) shows the integrated melted mass, in tons per Earth day, as a function of the lunar cycle for the JFCs (blue), the MBAs (green), the HTCs (purple), the OCCs (orange), and the overall flux (black).
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Panel (a) represents the total mass input, in g cm−2 s−1, as a function of the selenographic latitude and over a lunar month. Panel (b) shows the integrated mass inputs, in tons per Earth day, as a function of the lunar cycle for the JFCs (blue), the MBAs (green), the HTCs (purple), the OCCs (orange), and the overall flux (black).
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Latitude dependence of internal OH for two cases of diffusion activation energies, (a) E  a (OH) = 0.6 eV and E  a (O) = 0.55 eV and (b) E  a (OH) = 0.7 eV and E  a (O) = 0.65 eV, at new moon phase after 1000 yr. Dashed lines designate the local time. Both expectedly exhibit a latitude dependence, with the colder higher latitudes retaining more OH as a consequence of the temperature dependence on diffusion and RD. As shown in panel (a), the lower activation energy of diffusion results in an overall quicker dissipation of the implanted solar wind protons via transport to the surface and finally loss through RD, resulting in less hydroxyls. Conversely, the higher activation energy of diffusion slows diffusion from the subsurface to the surface, allowing more retention of OH throughout.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Snapshots of water formation for RD (top) and IM (bottom) assuming diffusion activation energy E  a (OH) = 0.6 eV and E  a (O) = 0.55 eV at various times. Initially, the hydroxyl defects from solar wind implantation are trapped within the immediate subsurface, reflected in the dominant production of water in the equatorial regions of the IM channel. Due to lower temperatures and lower solar wind flux at the higher latitudes, more time is required to build up the necessary density of surficial hydroxyl sites, resulting in a slow broadening of water formation in the RD channel localized within the warmer midlatitudes. Over time, regions in the high colder latitudes continue to accumulate solar wind protons, resulting in an apparent migration of peak water production in the IM channel toward the poles.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Variability of mass influxes of meteoroid impacts, in g cm−2 s−1, as a function of the selenographic longitude and latitude determined by the MIF for JFCs (panels (a) and (b)), ASTs (panels (c) and (d)), the HTCs (panels (e) and (f)), and the OOCs (panels (g) and (h)), in two Moon phases: first quarter (Day 8; left panels) and full moon (Day 15; right panels). Note that the color scales are different for each meteoroid population, and the corresponding fluxes are plotted over an elevation map of the Moon centered on the prime meridian. The vertical black lines represent the positions of the antihelion (LT = 0 hr, nightside), apex (LT = 6 hr, dawn terminator), and helion (LT = 12 hr, noon side) directions.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Adsorbed water concentration after 1000 yr for case 1 of diffusion activation energies discussed above. Dayside terminators are marked with black vertical lines. The left panel displays a clear diurnal variation of ∼30 ppm at the equator in adsorbed water. The right panel shows the total amount of water delivered to the PSRs. Note that this wt% is scaled assuming that the water is adsorbed onto the top layer of regolith grains only. Impact gardening will deliver fresh regolith to the PSRs within this time frame and significantly dilute the overall wt%.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Comparison of derived water concentrations from Honniball et al. (2020) with simulations (black line) as a function of latitude for various times of day on the lunar surface. As shown here, the general trend of our predictions is consistent with observations. However, the magnitude for early afternoon is generally ∼2× less than that observed, whereas the magnitude in difference between observation and prediction increases to ∼4–6× as dusk approaches and the regolith cools. The increase in the difference between model and observations at later times may be related to the observational data calibration. Alternatively, it may suggest a relatively large source of water generated and/or deposited on the nightside, possibly from IM.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Comparison of diurnal variation in lunar hydration as reflected in the change of spectral slope at 160 nm from Hendrix et al. (2019) with our simulated data. Our simulations predict a slightly higher (3×) coverage of water than implied by the LAMP data. In addition, there is no evidence in the LAMP data of the pre-dawn enhancement from molecules bouncing back to the nightside as predicted here and by Schorghofer (2015), which is possibly observed in the data set of Honniball et al. (2020).
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Comparison of lunar soil hydration levels derived from ground-based telescopic observations by Honniball et al. (2020) to simulated water concentration. Data from observations are shown in red, while data from simulations are shown in black, with the scale for the latter on the right and the scale for the former on the left. As seen here, the overall trend is similar, with hydration levels inversely proportional to the local temperature. However, the magnitude of the difference in observations is approximately twice what simulations predict. The peak near dawn may be related to the snowplow effect, as desorbing molecules hop immediately back to the cold nightside.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Shown here is the time of day dependence on the implanted hydroxyl at several latitudes. Examination of the pre-dawn (6 hr) and post-midnight (18 hr) reveals a ghost feature in the hydroxyl band across all latitudes due to shielding from Earth’s magnetotail. In addition, of note is the slight offset of minimum from noon. This delay arises from the necessary buildup of surficial OH sites before RD reactions are started and the continued replenishment from the cache of subsurface hydroxyls.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Comparison of (a) simulated internal hydroxyl, (b) adsorbed water, and (c) the summation of the two as a function of latitude for different lunar time of day to the reduced M3 data from Li & Milliken (2017). This figure highlights the similarity in latitude profile between observations and our model of solar-wind-generated OH along with the disparity of adsorbed water and observations. The discrepancy of adsorbed water compared to observations is a consequence of experimental binding energies of water and suggests that observations of M3 are a mixture of adsorbed water and internal hydroxyl species.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Average residual molecular water trapped within the regolith grain for the two cases of diffusion activation energies (a) E  a (OH) = 0.6 eV and E  a (O) = 0.55 eV and (b) E  a (OH) = 0.7 eV and E  a (O) = 0.65 eV. As shown here, the average is extremely low. However, actual impact events are grain to grain and thus would convert the entire concentration of hydroxyls to water with some remaining trapped. This would result in ∼0.5% of the regolith grains with roughly 100× more water than shown here. Our results then agree with observations of Honniball et al. (2021) assuming that SOFIA is sensitive to this relative concentration of grains with high water content.
Paper: Water Generation on the Moon from Solar Wind and Meteoroid Impacts
Authors: Brant M. Jones, Juan Diego Carrillo-Sánchez, Diego Janches, Menelaos Sarantos, Thomas M. Orlando
Figure caption: Shown here is the average water concentration of the PSR in units of ppm as a function of time. Initially, water does not reach the PSRs, as there are still available binding sites for adsorption elsewhere on the lunar surface. Over time, the majority of these sites are filled in and water must traverse larger distances to find a stable binding site. Naturally, these are the colder latitudes toward the poles. As water accumulates in the colder poles, there is a higher probability of landing in a PSR, where it then resides indefinitely.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: Phase changes of volatiles during their transport throughout the PSN. Volatiles are initially delivered as pure condensates in the PSN. Case 1 and Case 2 correspond to the scenarios where solids can only crystallize as pure condensates and can crystallize as pure condensates and clathrates, respectively. Pure condensate and clathrate pebbles are represented in blue and brown colors for a single species, respectively. Vapor is represented as purple circles. The condensation and clathration lines are represented as blue and brown solid lines, respectively. Once delivered to the disk, the phase of a given species is determined by the relative positions of the corresponding condensation, hydration, or clathration lines. Except for CO2, which vaporizes at a higher temperature than its associated clathrate, hydration, or clathration lines of the considered volatiles are closer to the Sun than their respective ice lines. Gaseous volatiles condense or become entrapped (depending on the availability of water ice) when diffusing outward of the locations of their condensation, hydration, or clathration lines. Conversely, volatiles condensed or entrapped in grains or pebbles are released in vapor form when drifting inward of their lines.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: Radial profiles of the C/H ratio relative to its protosolar abundance, calculated as a function of time in the PSN and in Case 2, assuming α = 10−3. Top and bottom panels correspond to the C/H ratio calculated in pebbles and in pebbles + vapors, respectively. Vertical dashed lines indicate CO ice line locations across PSN evolution (each color corresponds to a different epoch).
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: Molar abundance profiles of various volatile species in gaseous (orange lines) and pure condensate (blue lines) forms, calculated at different epochs (104, 105, and 106 yr) of the PSN evolution (Case 1). The species shown are H2O, NH3, CO2, H2S, Xe, CH4, PH3, Kr, CO, Ar, and N2.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: (a) Temporal evolution of elemental ratios (Ar/C, N/C, O/C, Kr/C, P/C, S/C, and Xe/C) calculated in Case 1 for solid condensates at the CO ice line within the PSN. The ratios are adjusted to their protosolar values. The P/C, S/C, and Xe/C ratios appear superimposed in the panel (see text for details). The Ar/C ratio is not represented in the solids because it falls below 10−25 (see text). (b) Same elemental ratios as in panel (a), but summed over both solid and vapor phases within the PSN. (c) Temporal evolution of the position of the CO ice line within the PSN.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: Molar abundance profiles of various volatile species in gaseous (orange lines), clathrate (green lines), and pure condensate (blue lines) forms, calculated at different epochs (104, 105, and 106 yr) of the PSN evolution (Case 2). The species shown are H2O, NH3, CO2, H2S, Xe, CH4, PH3, Kr, CO, Ar, and N2.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: From top to bottom: profiles of the disk midplane temperature, pressure, and surface density calculated at t = 104, 105, and 106 yr as a function of heliocentric distance, assuming α = 10−3.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: Same as Figure 6, but for Case 2.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: Top panel: elemental enrichments with respect to protosolar abundances predicted in the envelope of Uranus in Case 1, assuming a C enrichment ranging between 40 and 50 times the protosolar value. Dashed and solid lines correspond to predictions of the lower and upper values for the elemental enrichments, respectively. Bottom panel: same as the top panel, but for Neptune, assuming a C enrichment ranging between 68 and 104 times the protosolar value.
Paper: Insights on the Formation Conditions of Uranus and Neptune from Their Deep Elemental Compositions
Authors: Olivier Mousis, Antoine Schneeberger, Thibault Cavalié, Kathleen E. Mandt, Artyom Aguichine, Jonathan I. Lunine, Tom Benest Couzinou, Vincent Hue, Raphaël Moreno
Figure caption: Same as Figure 8, but for Case 2.
Paper: The Stability of a Dense Crust Situated on Small Bodies
Authors: Yoshinori Miyazaki, David J. Stevenson
Figure caption: (A) The effective viscosity of ice I adopted from Goldsby & Kohlstedt (2001) and Bland et al. (2009) as a function of temperature T for diffusion creep (dotted) and grain boundary sliding (GBS; solid) under a shear stress of 0.1 MPa. Three different grain sizes (0.01 mm, purple; 1 mm, orange; 10 cm, red) are considered here. (B)–(D) The instability timescale of a dirty-ice layer as a function of thickness h for g = 1.4 m s−2 and Δρ = 100 kg m−3. Timescale is calculated for three surface heat flux values: 80 (red), 20 (orange), and 4 mW m−2 (purple). Viscosity is estimated based on two different grain sizes ((B) 1 cm; (C)–(D) 0.1 mm) and a shear stress of 0.5 MPa, and shaded areas delineate when the shear stress varies between 0.05 and 5 MPa. The surface temperature assumed in (D) (T  s  = 90 K) is lower than in (B) and (C) (T  s  = 130 K). The age of the solar system (4.5 Gyr) is shown as a reference (gray). Whereas the decreasing trend in timescale for small h values can be attributed to the explicit effect of the thickening crust h, its rapid decrease for larger h values is due to the temperature-induced reduction in η 0.
Paper: The Stability of a Dense Crust Situated on Small Bodies
Authors: Yoshinori Miyazaki, David J. Stevenson
Figure caption: (A), (B) Contours of the instability timescale as a function of dirty-ice layer thickness and the surface heat flux for two different surface temperatures: (A) 130 K, relevant to Ganymede and Callisto (g = 1.4 m s−2), and (B) 150 K, relevant to Ceres (g = 0.27 m s−2). A grain size of 1 mm, a background stress level of 0.5 MPa, and Δρ = 100 kg m−3 are assumed. A dashed line separates the two different regimes after the onset of the RT instability (τ th ≶ τ RT). In the upper-right region (τ th ≫ τ RT), repeating instabilities are possible, resulting in the destabilization of the entire crust, whereas in the gray region (τ th < τ RT), the newly formed sublayer rapidly cools, and the instability only delaminates a layer of thickness L. Toward the lower-left corner, the viscosity scale height and the crustal thickness become comparable (dotted lines of h = 2L are shown as a reference), and the instability should result in some resurfacing. (C), (D) The viscosity scale height L and the critical length L crit, defined in Equation (10), are plotted for three different surface heat fluxes: 80 (red), 20 (orange), and 4 mW m−2 (purple). Shaded areas delineate when the shear stress varies between 0.05 and 5 MPa, with the middle line representing 0.5 MPa, and a grain size of 1 mm is assumed.
Paper: The Stability of a Dense Crust Situated on Small Bodies
Authors: Yoshinori Miyazaki, David J. Stevenson
Figure caption: Schematic illustration describing the aftermath of the RT instability, distinguished by the respective magnitudes of diffusion (τ th) and instability timescales (τ RT). (A) For τ th < τ RT, thermal diffusion swiftly reestablishes the initial thermal structure prior to the onset of the instability. If a phase transition accounts for the crust–mantle density difference, as seen in Io’s magma ocean, thermal diffusion will reinstate the thickness of the dense crust, leading to the repeating instabilities at the same depth. These instabilities will be contained in the thin layer at the bottom and do not destabilize the entire crust. (B) For τ th ≫ τ RT, the instability can propagate to the shallower regions, potentially destabilizing the entire crust. Temperature is illustrated using colors (blue, cold; red, hot), and regions occupied by lighter materials are represented with dots.
Paper: The Stability of a Dense Crust Situated on Small Bodies
Authors: Yoshinori Miyazaki, David J. Stevenson
Figure caption: (A) The temperature profile within the crust of Io (solid) and the viscosity scale height (L, dashed), assuming a crustal thickness (h) of 30 km and a crustal subsidence velocity of v  z  = 2 cm yr−1. Colors indicate different temperatures assumed at the base of the crust, but their differences are insignificant. (B) The growth timescale of the RT instability (1/ψ) as a function of the base temperature of the crust. Three different grain sizes are assumed: 0.1 mm (red), 1 mm (orange), and 1 cm (purple). The instability timescale (τ RT) increases with the square of the grain size but does not depend on the crustal thickness h. The diffusion timescale (τ th) needed to restore the thermal structure after the onset of the instability is also shown (black dashed), which is shorter than τ RT for a wide range of parameters.
Paper: The Stability of a Dense Crust Situated on Small Bodies
Authors: Yoshinori Miyazaki, David J. Stevenson
Figure caption: (A), (B) Dispersion relations of the RT instability for a denser elastic layer overlying a lighter one. (A) Colors indicate different layer thicknesses (10 km, green; 102 km, orange; 103 km, red), and the black line corresponds to when both layers are infinitely thick (Terrones 2005). Shear moduli of 105 and 104 Pa are assumed for the upper and lower layers, respectively, and the density difference is set to Δρ = 100 kg m−3. For results with finite-thickness, rigid boundary conditions (dotted) and free surface (solid) are assumed, but the maximum growth rates show little difference between the two conditions. (B) The same as (A) but varying the values of μ 1 (105 Pa, black; 107 Pa, red; 109 Pa, orange). The results for Δρ = 10 kg m−3 (dotted) and μ 2 = 102 Pa (dashed) are also plotted. (C) The minimum (critical) wavelength to induce the RT instability as a function of μ 1. Two values of Δρ (10, dotted; 100 kg m−3, solid) are shown.
Paper: The Stability of a Dense Crust Situated on Small Bodies
Authors: Yoshinori Miyazaki, David J. Stevenson
Figure caption: Dispersion relations of the RT instability for a denser viscous layer overlying a lighter one. (A) Nondimensionalized growth rate (ψ 0 in Equation (7)) as a function of dimensionless wavelength (kL). Three different viscosity profiles are considered: exponential, ﹩\eta ={\eta }_{0}\exp (-z/L)﹩, solid; linear, η = η 0(1 − z/L), dashed; and Arrhenius, ﹩\eta ={\eta }_{0}\exp (E/R(1/T-1/{T}_{0}))﹩, dotted. All three cases have the same viscosity derivative at the interface (η/(d η/dz) = −L). For the Arrhenius profile, a linear temperature increase is postulated, leading to a steeper viscosity increase than an exponential profile. (B) The same as (A) but plotting results of exponential viscosity profiles with dimensions for η 0 = 1016 Pa s, Δρ = 1000 kg m−3, and g = 1.4 m s−2. Colors indicate different viscosity scale height thicknesses (10 km, red; 105 km, blue), and a case for constant viscosity is also shown (green; Chandrasekhar 1961). Results for different viscosity profiles are also plotted for the case of L = 105 km, showing that exponential and Arrhenius profiles exhibit nearly identical growth rates under the same viscosity derivative.