xrp 2015 origin of the volatile envelopes of small...

Origin of the volatile envelopes of small-radius exoplanets.

For submission to NASA Exoplanet Research Program 2015 (NNH15ZDA001N-XRP)

1. Table of contents. .......................................................................................................... 0 2. Scientific/Technical/Management. .............................................................................. 1 2.1 Summary. .................................................................................................................. 1 2.2 Goals of the Proposed Study. .................................................................................... 2 2.3 Scientific Background ............................................................................................... 2 2.3.1. Kepler planets 1.6 R⊕ < R ≲ 2.5 R⊕ typically have O(0.1–1 wt %) volatile envelopes whose source and timing is unknown. ................................................. 3 2.3.2. Outgassing has well-understood physical and chemical controls that have not previously been applied to exoplanets. ........................................................... 5 2.3.3. Combining outgassing with the modern understanding of dynamics can constrain the source of the volatile envelopes on Kepler’s small-radius planets. 6 2.4 Technical Approach and Methodology. .................................................................. 6 2.4.1. Description of geochemical model and dynamical model. ................................. 7 2.4.2. Simulation of outgassing during embryo growth. ............................................... 9 2.4.3. Simulation of outgassing on young planetary surfaces T < 900 K. .................... 11 2.4.4. Simulation of outgassing by water delivery to planetary surfaces T > 900 K. ... 12 2.4.5. Confronting the outgassing hypothesis with Kepler (+ K2 + TESS) data. ........ 13 2.4.6. Assumptions, caveats, and refinements. ............................................................. 14 2.5 Perceived Impact of Proposed Work. ................................................................... 14 2.6 Relevance of Proposed Work. ................................................................................ 15 2.7 Work Plan. .............................................................................................................. 15 2.8 Personnel and Qualifications. ................................................................................ 15 3. References. ................................................................................................................... 16 4. Biographical sketches. ................................................................................................. 27 5. Summary of Work Effort. .......................................................................................... 22 6. Current and Pending Support. ................................................................................... 33 7. Budget Justification. .................................................................................................... 29 7.1. Budget narrative. ................................................................................................... 36 7.1.1. Facilities and Equipment. ................................................................................... 36 7.2. Budget Details. ........................................................................................................ 36 7.3. Data Management Plan. ......................................................................................... 36 8. Subcontract to The Pennsylvania State University. ................................................ 39

Origin of the volatile envelopes of small-radius exoplanets

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2. Scientific/Technical/Management.

2.1 Summary. Understanding the origin of the volatiles is fundamental to the interpretation of discoveries from the Kepler, K2 and TESS missions, as well as future observations of planets and planet-forming disks with the James Webb Space Telescope. Measurements of densities and masses for Kepler planets indicate that many small-radius planets contain significant volumes of volatiles. However, the relative contributions of water versus nebular-disk hydrogen to these planets’ volatiles inventories are not understood for most small Kepler planets. Current observations are consistent with either hydrogen accretion from the nebular disk onto rock cores, or hydrogen production via iron oxidation at water-rock interfaces (outgassing). These two divergent scenarios have very different implications for planet formation, evolution, and habitability.

We will test the H2-production hypothesis by combining basic models of planet thermal evolution and outgassing with a modern dynamical model incorporating disk migration. Specifically, we will model the rock-volatile interfaces of rocky cores as they grow and migrate, accounting for water-rock reactions, surface-interior recycling, volatile partitioning between magma and atmosphere, and hydrogen loss to space, both before and after the dispersal of the gas disk. We will consider two scenarios: (1) a “maximal production” scenario in which water is delivered to the cores in quantities that maximize hydrogen production (and thus planet radii), (2) a “dynamical” scenario in which the supply of water is set by dynamical models for radial mixing of small bodies that are enriched in water due to the location of their condensation relative to the ice line. In our “dynamical” scenario, we will explore a two-dimensional parameter space varying the total disk mass in solids, and the convexity of the solid body growth rate as a function of orbital distance (which affects radial drift/orbital migration). In both scenario (1) and scenario (2) we will self-consistently track pressure, temperature, and rock composition at the volatile-rock interface and their effect on H2 production, and set accretion of nebular H2 to zero. Thus, we shall obtain both: (1) strict upper limits, and (2) realistic estimates, for the radii of planets with atmospheres produced primarily via water-rock reactions.

We will compare the results of our calculations to constraints for the abundance of volatiles for small planets discovered by Kepler, K2, and eventually TESS, incorporating mass constraints from Radial Velocity (RV) observations and analyses of Transit-Timing Variations (TTV). For planets with radii measured to exceed our calculated radii in the “maximal production” scenario, we will obtain a lower bound on the contribution of nebular accretion to the volatile inventory of those planets. These results will lead to constraints for disk lifetimes, disk dynamics and astrophysical-metal enrichments of those planets (at their birth annuli). For planets with radii measured to be consistent with our model, we will predict their hydrogen:water ratios. Planets for which our model predicts large hydrogen:water ratios would not have Earth-like atmospheres, irrespective of our assumption of minimal nebular accretion. Therefore, our model will offer a valuable tool for assessing which planets could plausibly have an Earth-like atmosphere. Applying our model to small planets that lack mass and density measurements (e.g., in or near the habitable zone of Sunlike stars and thus beyond the reach of current RV or TTV) would inform the efficient prioritization of targets for follow-up observations, such as intensive RV campaigns or observations by JWST. The proposed work is relevant to the XRP call because it will improve our understanding of the chemical and physical processes of exoplanets (including their surfaces, interiors, and atmospheres) as well as the origins of their volatile envelopes.


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2.2 Goal of the proposed study. The objective of the proposed work is to constrain the role of outgassing (H2 production as a consequence of water-rock interaction) in shaping the volatile envelopes that are required to explain observed masses and radii of Kepler exoplanets 1.6 R⊕< R ≲ 2.5 R⊕, the most common planets around Sunlike stars (Fressin et al. 2013, Silburt et al. 2015). In order to test the hypothesis that outgassing can make a major (>50%) contribution to most (>50%) of these volatile envelopes, we will find the maximum outgassed H2 that can be produced and retained consistent with known chemistry and physics, and confront the results with the minimum H2 envelopes required by data. Achieving this objective involves the following steps (each including dynamical forcings and an escape-to-space parameterization):

Step 1. Simulation of outgassing during embryo growth (§2.4.2). Step 2. Simulation of outgassing on young planetary surfaces T < 900 K (§2.4.3). Step 3. Simulation of outgassing by water delivery to planetary surfaces T > 900 K (§2.4.4). Step 4. Confronting the water-rock hypothesis with Kepler (+ K2 + TESS) data (§2.4.5).

In order to define a focused, well-posed investigation of appropriate scope for a three-year study, we make several simplifying assumptions, which are explained and justified in §2.4.6.

2.3 Scientific background.

Rapid condensation of solar-composition gas at 140K (~3 AU) forms comparable masses of metallic iron, Mg-silicates, and H2O, but no iron oxides and no H2 (Lewis 2004, Krot et al. 2000). Instead, H2 production occurs secondarily on planetesimals (and larger bodies) via the oxidation of metallic Fe

Fe0 + H2O à Fe2+O + H2 (Reaction 1, “iron oxidation”)

(e.g. Rosenberg et al. 2001, Lange & Ahrens 1984, Dreibus & Wanke 1987); and during hydration of Fe-silicates on rocky bodies, by reactions such as:

6Fe2SiO4 + 11H2O à 2Fe3O4 + 5H2 + 3Fe2Si2O5(OH)4 (Reaction 2, “serpentinization”)

which in a planetary context will always occur alongside hydration of Mg-silicates (Fruh-Green et al. 2004, McCollom & Bach 2009, Klein et al. 2013). Reactions like (2) are most likely at the volatile-rock interface (“surface”) of a growing planet. When Tvri exceeds ~900K, Reaction (2) is thermodynamically unfavorable (e.g. Abe & Matsui 1986). H2 production can continue, if accreted objects contain both incompletely oxidized iron and H2O (Reaction 1). Volcanism can contribute to outgassed atmospheres, and is important for moderating climate over Gyr (e.g. Kite et al. 2011); but the volcanic flux is too small (and likely also too high molecular weight) to explain the voluminous volatile envelopes that are the focus of this study (Kite et al. 2009).

It is reasonable to assume that exoplanets’ rocky cores have Fe contents similar to that of the Solar System’s rocky planets (~32wt% Fe by weight)1. Early in nebular history, this Fe would be in the Fe0 state. It follows that the hydrogen-producing potential of exoplanetary rocky cores is ~32 wt% × (3 e– / 2e– ) × ( 2 amu / 56 amu ) ≈ 1.7 wt% (Elkins-Tanton & Seager 2008a, 2008b, Rogers et al. 2011) – the “stoichiometric limit,” corresponding to a planet with all Fe in the Fe3+ state , and lacking a Fe-metal core. We will calculate the extent to which this H2-accumulation potential is tapped using well-characterized thermodynamics combined with simple assumptions

1 Only minor Fe/Si variations are expected between planet-forming disks (consistent with data; Jura & Young 2014, Dressing et al. 2015); varying Fe/Si will be considered in Year 3 (§2.4.6).


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about rock recycling and escape-to-space, forced by simulations of migration, accretion and impacts.

2.3.1. Kepler planets 1.6 R⊕ < R < 2.5 R⊕ typically have O(0.1–1 wt%) volatiles whose source and timing is unknown.

Fig 1. (a) (From Marcy et al. 2014) Occurrence rate of exoplanets. We hypothesize that the “plunge” in occurrence rate at ~2.5 R⊕ (Lissauer 2014) corresponds to the maximum hydrogen that can be outgassed

and retained during assembly of a ~6 M⊕ rock core (Rogers 2015). (b) (From Rogers et al. 2011) Maximum radii for exoplanets with outgassed atmospheres (pink lines), showing radii can double relative to the no-degassing case (black lines) from the outgassing of ~1 wt% H2. Blue and green lines show correspond to water-limited Fe-oxidation, which will be tracked in the proposed work. Losses are not considered in this plot, but will be included in our proposed work.

The observational target for our modeling is clear (Lissauer et al. 2012, Rowe et al. 2014). Observations of short-period planets require low-density envelopes to explain the densities of R > 1.6 R⊕ Kepler planets (Rogers 2015); smaller planets are consistent with Earthlike composition (Dressing & Charbonneau 2015), but could also have significant atmospheres. Although there is no proof that all of the envelopes must contain H2, many of these envelopes are thought to be H2-rich for the following reasons. (1) Simplicity; some of the envelopes require H2, and there is no requirement to invoke the extra ingredient of a massive ice envelope (e.g. Jontof-Hutter et al. 2014, 2015; Rogers 2015). (2) Cosmochemistry; worlds with a steam atmosphere but no H/He envelope are unlikely on cosmogonic grounds (Nettelmann et al. 2011, Valencia et al. 2013). (3) Evidence for photoevaporative sculpting of the Kepler ensemble; this suggests that short-period water-worlds are uncommon, because evaporative sculpting is ineffective for water shrouds though effective for hydrogen shrouds (e.g. Youdin 2011, Dong & Zhu 2013, Owen & Wu 2013, Fortney et al. 2013, Wu & Lithwick 2013, Lopez & Fortney, 2013). Assuming a H2-rich composition, envelopes typically comprise ~0.1-1 wt% of the core mass (Saumon et al. 1995, Rogers et al. 2011, Lopez & Fortney 2014, Wolfgang & Lopez 2014, Silburt et al. 2015, Howe et al. 2014). These low envelope-mass values would drop ~10x further if a H-dominated He-free atmosphere (as appropriate for outgassed envelopes) were assumed instead of the H/He mix (as appropriate for nebular accretion) that is used in estimates published to date. Low inferred envelope masses underline the stoichiometric sufficiency of iron oxidation, but pose an intriguing challenge to the runaway accumulation of H2 predicted by core formation theory. To reconcile core formation theory with nebular accretion of these hydrogen envelopes, fine-tuning is required (Lee et al. 2014, Bodenheimer & Lissauer 2014, Inamdar & Schlicting 2015). But fine-tuning is hard to reconcile with reduction in planet occurrence rate around 2.5 R⊕ (Petigura et al. 2013, Siliburt et al. 2015), which requires that an account of Super-Earth volatile-envelope origin via


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nebular accretion must explain not only how nebular accretion can stop at the right size, but also why this stopping is generic. This is not easy (e.g., Jin et al. 2014). While it is possible that nebular accretion may be brought into line with observations (e.g. Ogihara et al. 2015a, 2015b), these tensions with data are intriguing, and motivate consideration of alternative hypotheses such as outgassing. Outgassing can naturally explain “stopping” at O(0.1-1) wt% H2, because of the stoichiometric limit: there are only enough Fe atoms to make ~2 wt% H2 envelopes (Rogers et al. 2011). If outgassing survives a detailed test, then there might be major implications for the extent of volatile-rock interactions and the structure of small exoplanets (Elkins-Tanton & Seager 2008b). Despite its potentially important implications, the outgassing hypothesis has never been tested for R < 2.5 R⊕ planets; it is this gap that the proposed work will fill.

___________________________________________ Fig. 2. Theory favors inward migration of Super- Earths during their formation, favoring accumulation of volatiles whose reaction with Fe0 would outgas H2 (After Raymond, KITP Exoplanets Meeting, 2015).

To test the outgassing hypothesis, we must determine to what extent the H2-producing capability defined by the “stoichiometric limit” is tapped, and to what extent the H2 so produced can be retained during accretion to contribute to the observed envelopes.

This requires modeling both the timing and temperature-pressure (T-P) conditions of water-rock interaction during planet formation, as well as the pace and timing of water delivery. Modern dynamics implies stochastic variations both in water delivery and in giant-impact loss, though with orderly trends related to distance from the ice-line. Kepler’s systems of small tightly-packed inner planets (STIPs) have orbital architectures and low eccentricities that suggest limited excitation by dynamical instabilities, and orderly within-disk migration (Fabrycky et al. 2014). Migration is probably required to concentrate solids in the innermost part of the disk (Raymond & Cossou 2014, Schlichting et al. 2014, Schlaufmann 2014, Bodenheimer & Lissauer 2014, Ogihara et al. 2015a; see also Chiang & Laughlin 2013). Migration is consistent with the outgassing hypothesis for the envelopes of short-period small-radius planets (Fig. 2).

2.3.2. Outgassing has well-understood physical and chemical controls that have not previously been applied to exoplanets.

The physical and chemical controls on Reactions (1) and (2) are fairly well understood (e.g., Zolensky et al. 1989, Zolensky et al. 2008). This is partly because massive H2 outgassing via Reactions (1) and (2) is required to explain Solar System planet-composition and meteorite data. It has been known for decades that solar system rocky planet geochemistry requires heterogenous accretion with significant variation in the redox state of accreting materials due primarily to varying levels of iron-water reaction. Now, these data can be matched by dynamical simulations (Raymond et al. 2014, Rubie et al. 2015). These modern dynamical simulations make a compelling case for the importance of radial mixing and stochastic late-stage accretion in volatile delivery, and cap six decades of research into Solar System volatile origins (e.g. Atreya et al. 1989, Porcelli et al. 2002). For example, the solar system’s rocky objects are generally oxidized by isochemical water-rock interaction, with little evidence for nebular accretion (Wasserburg 1964, Brearley 2006, Bland et al. 2009, Halliday 2014; but see also Pepin 2006). In particular,


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~15% of Earth’s mass must derive from accretion of objects that were oxidized by Reaction (1) in order to explain the Fe2+ in the Earth’s mantle (Ringwood 1979, Dreibus and Wanke 1989, Schönbächler et al. 2010, Rubie et al. 2015). In the Solar System, oxidation of Fe0 commonly occurred on small objects (GMm / kTR < 1) so that H2 was lost to space (Young 2001). This reduced the H2-producing potential of later stages; atmospheres outgassed later in accretion were high-molecular-weight (Schaefer & Fegley 2009). Removal of hydrogen from young solar system objects by thermal escape and impact erosion (Zahnle & Kasting 1986) was also efficient, and the Solar System’s rocky planets (R < 1R⊕) lack hydrogen. Intriguingly, giant impacts did not remove all the volatiles, which suggests late volatile accretion and/or protection by dissolution into silicate melt (e.g. Mizuno et al. 1980).

Fig. 3. We propose to track H2 production via Fe-oxidation on planetary embryos (leading to irreversible loss to space) (§2.4.2), H2 production via serpentinization on young planets (in competition with thermal and nonthermal escape) (§2.4.3), and H2 production via Fe-H2O reaction near

magma-ocean surfaces (in competition with thermal escape, nonthermal escape, and giant-impact atmospheric ablation) (§2.4.4). The predicted H2 envelopes will then be compared to Kepler data (§2.4.5).

The controls on the efficiency of H2 production by serpentinization (e.g., Reaction 2) have been intensively studied both empirically and theoretically (Sleep et al., 2004; Kelley et al., 2005, McCollom & Bach 2009, Klein et al. 2012, Klein et al. 2013). Rapid hydrogen production occurs when liquid water oxidizes Fe2+ to Fe3+ (or Fe0 to Fe2+) provided that a thermodynamic driving force is present, that water can interact with the bulk of the rock due to high permeability or tectonic resurfacing (Young et al. 2001), and that the Mg:Fe ratio permits H2 production.

In summary, contrast with Solar System data (R ≤ R⊕) implies that retention of outgassed hydrogen on R < 2.5 R⊕ planets would require some combination of 1) late reaction of water with reduced iron (cold accretion followed by melting on planetary embryos), 2) retention of hydrogen within planetary interiors to protect it from giant impacts, 3) attenuated escape-to-space, or 4) late delivery of volatile-bearing objects. Many of these processes would be assisted by the large surface gravity of super-Earths, because both accretion and loss are sensitive to surface gravity. However, all of these processes are missing from detailed models of planetary volatile evolution (Fig. 4), and so these are the processes that we propose to investigate, as detailed in §2.4.

DYNAMICALLYCONSTRAINEDWATER SUPPLY

Time (yr)TARG

ET: VOLATILE EN

VELOPES O

N

KEPLER’S SMALL-RAD

IUS EXO

PLANETS (§2.4.5)

BASICTHERMODYNAMICMODELS TRACKH2 OUTGASSING metal core

may form

type Imigration

giant impacts

108 109105 106 107

accretedfrom

narrowannulus

stochasticdelivery

from largeseperations

resurfacing

serpentinizationwater delivery

well-mixedmagma ocean

dissolution

thermolysis

water delivery

heatingcracking

100

km -

2000

km

thermal escapenonthermal escapeimpact ablation

2000

- 16

000

km

2000

- 16

000

km

DYNAMICALLYCONSTRAINEDROCKY-COREGROWTH ANDMIGRATION

iron ox.1D grid

impermeableinterior

ice melts throughout interior

box model

ggiaiim

correormm

e Itioonn

1600

0 km

(§2.4.2) (§2.4.3) (§2.4.4)

rock surfaces melt

thermal escapenonthermal escape

thermal escape

accretion

porous interior


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2.3.3. Combining outgassing with orbital dynamics can constrain our understanding of the volatile envelopes on Kepler’s small-radius planets.

Fig. 4. Our proposal in the context of (a small subset of) published geochemical and dynamics models, showing that the proposed work is a natural extension of long-standing community-wide efforts towards using both geochemistry and dynamics to study the problem of volatile envelopes in planets interior to the ice line.

The formation history of Kepler's Systems with Tightly-packed, Inner Planets (STIPs) is still an area of active research. Some (e.g., Chiang & Laughlin 2013) have proposed that such systems may form in situ. In this scenario the

H2O abundance would be quite low, making outgassed atmospheres an implausible explanation for most of Kepler's planets 1.6-2.5 R⊕. This proposal aims to understand the maximum plausible atmosphere mass to be outgassed and survive to the present epoch. Therefore, we focus on models that involve significant orbital migration (Fig. 2), allowing STIPs to harbor small, inner planets that accreted significant H2O from planetesimals that formed beyond the ice line before drifting inwards due to gas drag (Fig. 5). While the gas disk dissipates random velocities efficiently, planet embryos will remain on nearly circular orbits, severely limiting close encounters. As the gas dissipates, planetary embryos will acquire larger eccentricities, leading to close encounters, collisions, scattering and radial mixing of material. This framework naturally explains the nearly packed nature of planetary systems, and the small fraction of systems in mean motion resonances. It also provides a practical means for coupling chemical and dynamical modeling. The roles of radial mixing, and net migration toward the star, are critical in setting pace and extent of volatile delivery. We propose to combine this new dynamical understanding with the mature thermodynamics of water-rock interaction.

In summary, the proposed work will bring together the knowledge base on water-rock interactions and on the dynamics of planetary-system formation. Specifically, the outgassing hypothesis can be tested with existing data by combining constraints from thermodynamics, the new dynamics (Raymond et al. 2009, Dauphas & Morbidelli 2014, Raymond et al. 2014), giant impact ablation, the lower limit on H2 escape set by thermal (Jeans + hydrodynamic) escape (Donahue 1986), and the need for migration. In addition the hypothesis must confront strong inferences about early solar system processes from well-established meteorite data.

2.4 Technical Approach and Methodology.

Our hypothesis test involves two scenarios: (i) To what extent can H2 be produced and retained assuming optimal delivery of water, but with realistic thermodynamics, realistic parameterized giant impacts, and realistic escape-to-space? (ii) As (i), but with realistic delivery of water.

We identify several potential bottleneck steps each of which could individually limit a contribution from outgassing to observed volatile envelopes and how we will address each in our study: (A) Delaying the water-rock reaction until the embryo stage (big enough to retain H2) requires keeping the growing object cool (by low 26Al/60Fe, and/or pebble accretion), and/or late


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delivery of water-rich objects. à Thermal model of growing embryos. (B) Rapid migration (and H2 accumulation) may limit H2 production via serpentinization by heating the volatile-rock interface. à Use dynamical model to drive the physical model of serpentinization under a growing atmosphere. (C) Giant-impact ablation may remove H2 produced at earlier stages. à Parameterized stochastic giant-impact model. (D) Water supply may limit H2 production à Dynamical model of water delivery. These bottlenecks range from kinetically-limited to supply-limited, and predict correspondingly different size dependencies, atmospheric-loss threshholds, and volatile-abundance trends with planet radius and period.

Table 1. Trades in Design of our Numerical Experiment. Arrows (à) show Year 3 refinements. Process or variable

Choice Reasoning / Rationale for Choice

Hold constant

Consider range

Parameterized

Basic, physically

consistent model

State-of-the-art m

odel

Water-rock thermodynamics

✓ Well-understood process. Precise, accurate thermodynamic databases are available. Kinetics unimportant for {P, T} range of interest.

Thermal evolution of embryos

✓ A growing embryo can only marginally stay cool enough to retain H2, so 1D model required.

Type I migration ✓ à Important, but poorly understood process. 1D disk model required. Rock-mantle convection /resurfacing

✓ The exponential decrease of rock viscosity with increasing temperature buffers the temperature of large rocky objects.

Solid-body growth rate

✓ ✓ Growth rate varied as a function of distance.

Initial H2O content ✓ Ice content external to the H2O ice-line is varied: {5, 10, 50} wt %. Nebular Fe/Si ✓ à Cosmochemical variations are unlikely to change H2 production by

>30% of total H2. Carbon content. ✓ à Initially omitted for simplicity; CO snowline more distant than H2O

snowline. Photoevaporation efficiency

✓ Envelope origin is our goal, so we make the simple assumption of constant photoevaporation efficiency when comparing to data.

Permeability and porosity

✓ Small planets cannot retain internal overpressures of H2, and large planets have negligible internal porosity.

We will model each bottleneck using a geochemical model forced by dynamics, and supported by simple parameterizations or 1D models of thermal evolution and atmospheric escape. Given the large ratio of computational expense of N-body accretion codes to geochemical models, we will run many geochemical models based on a relatively small number of N-body runs. The analytic model will still include a stochastic rate of giant impacts. For Scenario 1, we will consider bottlenecks (A)à(B)à(C) in series, but supply water “from infinity” to the extent that maximizes H2 production (This might physically correspond to a rain of small comets). Because the total amount of water so supplied will be <10 wt% of solid mass (Elkins-Tanton & Seager 2008a), the “extra” mass of water will not cause grossly unrealistic dynamics. For Scenario 2, we will consider bottlenecks (A)à(B)à(C) in series and implement (D) self-consistently.

Our proposed work will build on consensus ideas in geochemistry and dynamics (Fig. 4), and advance the state-of-the-art by combining these simple constraints (Fig. 4).


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2.4.1. Description of the geochemical model and the dynamical model. Separations of time-scales and length-scales define our modeling approach. Time: The giant impacts that heat planetary mantles and strip hydrogen occur on a much shorter timescale than the migration (and relatively slow accretion) during which hydrogen production occurs. We exploit this decoupling of timescales by running water-rock interaction simulations with a fixed flux of volatile input, and use the output from a grid of such runs to obtain a parameterization for forcing by N-body simulations of giant impacts and planet migration. Length: The embryo code needs to track entropy versus depth, but heating and possible core formation causes growing planets ~2000 km in radius to transition to an ~isentropic state. Therefore, we will write separate codes for the temperatures and reactions occurring in embryos and on planetary surfaces: a 1D code for embryos (radially symmetric; §2.4.2) and a parameterized code for well-stirred planets interacting with a volatile envelope that may or may not include a separated aqueous phase (§2.4.3, §2.4.4). The geochemical codes are computationally cheap and have negligible feedback on the N-body results, so we use the N-body code to drive the geochemical evolution of every body in the simulations (geochemistry is applied in postprocessing).

For the embryo-interior simulations we will use the CHIM-XPT geochemical code (Bethke 2010), which has been applied to Mars problems by the Postdoctoral Researcher (Melwani Daswani et al., accepted). We shall use the SUPCRT92 and Robie & Hemingway (1995) databases. The full-chemistry capabilities of CHIM-XPT are not used:- instead, chemistry is specified using a simple, fixed chemistry (Mg-Fe-Si-O-H), and all other elements (except water) are zeroed out, which allows all the most important reactions while simplifying the analysis. Initial Fe:Mg:Si ratios are set to match solar system values. Consistent with theory for water-ice condensation and snowline locations (Kok et al., 2012), we assume that the H2O-ice concentration at t=0 is zero inward of the snowline for our planetesimals and steps upwards moving outwards. The initial H2O content of planetesimals exterior to the iceline (Wil) is an important control on the ultimate Fe:H2O ratio of interior planets that are more detectable by Kepler. We shall consider (1) a “dry” endmember case (Wil = 5 wt% H2O); (2) a “Solar System analog” case, guided by the ~10 wt% H2O in the CM2 parent body (e.g., Clayton & Mayeda 1996). (3) Wil = 50 wt% H2O, guided by equilibrium nebular condensation chemistry (Ciesla et al. 2015). Initial and boundary conditions are chosen self-consistently so that large quantities of CO2 ice are not incorporated. (More sophisticated analyses with full chemistry support the assumption that H2 is the dominant gas for the large water abundances we assume; Schaefer & Fegley 2007). In order to be certain of covering the highest temperatures and pressures encountered on large hot rocky cores, we will supplement the databases used in CHIM-XPT database with the NIST Reference Fluid Thermodynamic and Transport Properties Database version 9.1 (REFPROP). This database will be purchased with Kite’s startup funds.

The dynamical model consists of planetesimals-embryo and embyro-embyro agglomeration via direct n-body integration with semi-analytic migration, and analytic embryo growth and injection at the outer boundary of the disk (~3 AU). The n-body evolution will be implemented using Mercury (Chambers 1999) and/or Rebound (Rein & Spiegel 2015), using a high-order adaptive timestep integrator (i.e., Burlisch-Stoer or IAS15) for the early phases and an integrator that preserves the symplecticity of Hamiltonian systems (i.e., hybrid MVS-BS or IAS15) for once the gas has dissipated and collisions are rare. The embryo growth will be implemented following Cossou et al. (2014). Migration will be modeled initially as a frictional damping term f = −f τdamp

−1 (Moorhead & Adams 2005) and could be easily extended to allow for separate migration


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and eccentricity/inclination damping timescales (Lee & Peale 2002). The gas will dissipate exponentially with a timescale of ~3 Myr (in the baseline model) until it has a negligible effect on both embryos and planetesimals (Rein et al. 2012). The baseline initial disk model will be based on a 5x the minimum-mass solar nebula and we will vary from 0.5 – 20x the minimum-mass solar nebular in 8 logarithmically-evenly-spaced steps. Each simulation will start with ~1-100 isolation mass embryos which interact with the star, disk, each other and a disk of ~1,000 "super-particles" that collectively represent the planetesimals. For the sake of computational speed planetesimal-planetesimal interactions and collisions will be neglected for the baseline calculations. We will augment the Mercury and/or Rebound codes to increase the sophistication of collisions to track total mass loss and mass loss of volatiles. Integrations will be for 108 yr.

From experience gained during prior N-body work by the Co-I (e.g. Ford & Rasio 2008) and the PI (e.g. Kite et al. 2015) using Mercury, ≤9 CPU months are required for 50-embryo 1000-planetesimal runs. The PI’s partition on UChicago’s Midway cluster easily satisfies all our CPU and storage requirements. Other initial and boundary conditions are varied as specified below. Based on the prior experience of both the PI and the Co-I with Mercury (Kite et al. 2015, Ford & Rasio 2008), time for analysis (and not CPU time), sets the length of the project. We set the number of runs, and the number of parameters to vary (Table 1), based on this prior experience.

2.4.2. Simulation of outgassing during embryo growth.

Any H2 produced when the ratio (GMm / kTexobaseRexobase) ≲ 2.5 will be lost to space via thermal escape (Volkov et al. 2011). This “early” iron oxidation represents an irreversible loss of water-reducing potential for later stages in planet formation, and must be tracked in order to self-consistently compare with data for Kepler planets. H2 production is unavoidable when Fe and H2O are inter-mixed at scales <10 m at T > 273K. T < 273K interiors are maintained during accretion for ~103 km-diameter objects provided that Short Lived Radioisotopes (SLRs) are absent (Solomon 1979, Toksoz et al. 1978, Hanks & Anderson 1969, McSween et al. 2002). Therefore, production of H2 on R<103 km bodies in the Solar System is due to SLRs. We posit that many or most planetary systems lacked the Solar System’s abundant SLRs, due to forming in a low-mass cluster or forming early in a massive cluster (Gaidos et al. 2009, Adams 2010, Gounelle 2014). Because in our Solar System SLRs heated small bodies to allow liquid water to react with iron, much H2 was produced by Fe-oxidation at the planetesimal stage of planet growth (Wilson et al. 1999, Brearley 2006, Rosenberg et al. 2001). This H2 was returned to the nebula by thermal escape (e.g., Guillou et al. 2015). We hypothesize that iron oxidation – and thus hydrogen production - can be delayed until surface gravity is high enough to allow the H2 to be retained. To test this hypothesis, we model the thermal evolution of nonmigrating planetesimals undergoing runaway growth (each with a H2O content appropriate to their birth location):

(Grimm & McSween 1989, Melosh 1990), where T(r,t) is temperature as a function of radius r, κ is thermal diffusivity, G is the gravitational constant, M is embryo mass, Fm is the accretion mass flux, R is embryo radius (R = 10 km at t = 0), σ is the Stefan constant, is is a boundary flag set to 1 at the boundary and 0 in the interior, Hc is chemogenic power, and Hr is radiogenic power. Fm (and thus R(t) ) is prescribed using Safranov-type accretion equations 5.33 – 5.36 in Armitage (2010), with runaway growth equations for R<10 km, oligarchic growth equations for R>100 km,


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and a logarithmic weighted mean between these endmembers at intermediate radii. We will compare the results from this baseline scenario with simple simulations meant to represent the streaming instability (e.g. Johanssen et al. 2015); in those cases, we won’t resolve individual collisions. Incoming thermal radiation is included, self-consistently with the degree of disk dispersal. Thermal evolution of 102 km < R ≲ 103 km objects in the Solar System has been extensively studied (e.g., Castillo-Rogez & McCord, 2010). We solve (1) using 100 grid points in radius (the grid expands to track growth). We assume runaway growth, for which impacting objects have radii much smaller than the growing embryo and deposit their energy in the surface layer, and we neglect shattering and erosive impacts. Regolith insulation and porosity is assumed to be small (in order for H2 to be retained against thermal escape, the planets will be large enough that the regolith contribution to internal temperatures will be small). We set Hc = 0 initially (Neveu et al. (2015) find that the thermal effect of silicate hydration is minor), but will carry out sensitivity tests to Hc ≠ 0 and Hr ≠ 0. The Hc = 0 case will be used when confronting our output with data (§2.4.5); this is conservative in terms of disproving the outgassing hypothesis.

Fig. 5. Example Mercury output. Each track corresponds to an embryo. Black asterisks correspond to giant impacts. The numbers in the right axis correspond to planet mass in Earths. Migration (Type I) continues until disk dispersal at 3 Myr. Similar calculations will be used to drive our geochemical code (although we adopt a narrower initial distribution of solid mass). Close-in super-Earths will have a significant volatile content in this scenario. (The output underlying this plot was generated by C. Cossou).

CHIM-XPT tracks H2 production via Reactions (1) and (2) where T>273K. We neglect kinetics (because Reactions (1) and (2) are strongly exothermic, leading to a positive feedback between reaction and reaction-rate near the site of the reaction; Grimm & McSween 1989). We assume temperatures are buffered to ≲300K by ice melting, consistent with chondrite paleothermometry (Guo & Eiler 2007, Rosenberg et al. 2001, Clayton & Mayeda 1999) as well as the output of more sophisticated reaction-transport models considering permeability feedbacks (e.g. Young 2001). This is reasonable because we are not interested in higher-temperature metamorphism for this proposal. Once produced internally, hydrogen is promptly delivered to the surface, because preexisting fractures are likely (Neveu et al. 2015) and H2-overpressure can from new fractures (Wilson et al. 1999). Thus, atmospheric thickness is thus set by balance between hydrogen production and thermal escape. We assume thermal escape with Texobase and Rexobase set self-consistently for a pure H2 atmosphere using standard methods (e.g. Pierrehumbert et al. 2010) assuming UV is absorbed high in the nebula.

The output of Step 1 will be a set of embryos in nested orbits, each with a self-consistent internal structure: bodies interior to the ~1.5AU will have negligible ice, and bodies exterior to ~1.5AU will be more-oxidized but may retain cool and ice-rich cores (a “jawbreaker” structure; Zahnle et al. 1988). These embryos will serve as the input to the next step (along with unresolved, unprocessed planetesimals of total density equal to that of the embryos).


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2.4.3. Simulation of outgassing on young planetary surfaces T < 900K We hypothesize that a simple model of serpentinization at the volatile-rock interface can set an upper limit on the contribution of serpentinization to exoplanet atmospheres.

Planets larger than 2000 km in radius will be hot throughout (e.g. Zahnle et al. 1988). Differentiation of a Fe-metal core may occur. This physical transition approximately corresponds to dynamical stirring and the onset of large-scale radial mixing (due to type I migration and scattering). In this Step, we are running an N-body simulation, and individual collisions are now resolved between the objects from §2.4.2.

The conditions on a young embryo may be ideal for serpentinization (Reaction 2 is a simplified representation of the full Mg-Fe-Si-O-H chemistry of serpentinization that will be calculated using CHIM-XPT). Serpentinization on Solar System habitable zone planets is sluggish because ocean-crust temperatures are low away from mid-ocean ridges. On young hot planets or on Super-Earths, these constraints are lifted: seafloor temperatures can easily reach 500 K (the optimal temperature for serpentinization), and rapid tectonic recycling can lift the supply limit. For example, serpentinization could theoretically have generated a 500-bar H2 atmosphere on Early Earth (Sleep 2004). However, as the H2 atmosphere thickens, temperatures at the base of the atmosphere rise toward the ~900K that would prevent further H2 production. We will track this important negative feedback by calculating surface temperature using a semi-gray analytic approximation to the radiative layer (Guillot et al. 2010) patched to a convective inner envelope (this compares well to more detailed calculations, e.g. Fortney et al. 2007).

Mass is supplied to the embryo via planetesimal accretion at a rate set by the dynamics code. (Embryo-embryo collisions trigger the magma ocean stage; §2.4.4). The water content of newly-accreted planetesimals is added to the volatile envelope, and the rock and metal content is mixed into the rock mantle. Provided that a water ocean exists, serpentinization proceeds very rapidly to thermodynamic equilibrium in a boundary layer with a thickness tuned to match the depth of penetration of hydrothermal cracking on Earth (~2 km; Vance et al. 2007, Alt et al. 1986, Neveu et al. 2015), with thickness scaled as (heat flow)-1 as appropriate for hotter planets (Kite et al. 2009). This surface boundary layer is added back to the well-stirred mantle (and replaced by average mantle material) at a rate set by standard parameterizations of mantle convection (water-saturated olivine rheology) with lid recycling (Kite et al. 2009, Schubert et al. 2001. Assuming lid recycling (Kite et al. 2009) helps to set an upper limit on H2 production, because stagnant-lid mantles limit the supply of fresh rock to the volatile-rock interface and thus will outgas less H2. Our model tracks rock-mantle Fe:Mg. H2 production from hydration of pure Fe-silicate is thermodynamically unfavorable, but hydration of pure Mg-silicate does not yield H2. Maximum yield occurs for Fe:Mg ≈ 1:1 (McCollom & Bach 2009). Mass is removed from the embryo by thermal escape (same parameterization as §2.4.2). When the Jeans parameter < 2.5, drag from escaping H2 also returns some H2O to space (Hunten et al. 1987). After gas-disk dispersal, EUV photoevaporation (assumed efficiency = 0.1; Lopez et al. 2012) occurs.

In our “optimal-degassing” scenario we assume that water is supplied to the planet in the optimal amount for further serpentinization (which requires maintaining an aqueous phase at the volatile-rock interface). In our “dynamically-realistic” scenario we take into account the solubility of H2O in H2 (following Rogers 2012) and the corresponding reduction of the thickness of the water ocean; in some cases, the water ocean will dry out. Dry-out does not prevent lid recycling (Kite et al. 2009, O’Reilly & Davies 1981), but we assume it prevents Reaction (2).


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The residence time of a growing planet in this Step may be brief. Some of the H2 produced during this “warm” stage will be eroded if giant impacts occur during the “hot” stage (§2.4.4), and some will be protected by dissolution into magma or favorable giant-impact energies. Some planets will (like Mars) evade both giant impacts and inward migration and never enter the high-temperature evolution modeled in §2.4.4, favoring retention of the H2 produced during this stage.

2.4.4. Simulation of outgassing by late delivery of water to planets with surfaces T > 900K. We hypothesize that Reaction 1 can make a major contribution to the low-density envelopes of Kepler planets (e.g. Hayashi et al. 1979, Sasaki & Nakazawa 1990, Harper & Jacobsen 1996, Hamano et al. 2013). As planets migrate towards the star, the thickening H2 cloak insulates the surface, and the era of giant impacts opens; a magma ocean will form. Magma oceans are well-stirred (Solomatov 2000), and we assume the volatile envelope equilibrates with the rock at each timestep. Serpentinization is now thermodynamically unfavorable, and H2 dissolves into the magma (Hirschmann et al. 2012, Mizuno et al. 1980, Porcelli et al. 2001). H2 solubility depends on the temperature at the volatile-magma interface, which we will continue to track (§2.4.3). Dissolution of the envelope into magma can protect part of the planet's H2 from loss during any given giant impact (Hirschmann et al. 2012). But the interior-surface re-equilibration time is short compared to the time between giant impacts, so exsolution will precede the next giant impact and allow that next impact to remove more H2.

Upon embryo-embryo or embryo-planet collision, Fe-metal cores are assumed to merge with ~50% efficiency; remaining Fe is dispersed through the rock mantle. This self-consistently maintains young planet rock-mantles that are more reducing than the iron-wustite buffer in agreement with Solar System geochemical data (e.g., Frost et al. 2008) and with core-formation models constrained by isotopic data (Nimmo & Agnor 2006, Rubie et al. 2015). We track Fe-metal-core size self-consistently. Si and O are not permitted to dissolve into the Fe-metal liquid. Volatiles (H2O, H2) are partitioned upon collision to the volatile envelope with 100% efficiency (although they promptly equilibrate with magma). The planet cools between giant impacts via radiation to space at the runaway-greenhouse limit.

The Saumon et al. (1995) equation of state for H will be used. This includes H2 dissociation, although we anticipate this will be minor.

The temperature dependence of the H2O(g) à H2 + 1/2O2 equilibrium constant is set to exp(6.347 – 29435/T) bar1/2 (Robie and Hemingway 1995, Krot et al. 2000). The solubilities of H2 and H2O in magma, and their temperature dependencies, are set to SH2 = exp(-13.662 + 1068/T) and SH2O = exp(-14.193 + 3890/T) where S has units of mol/m3/Pa (Hirschmann et al. 2012, Shackelford et al. 1972, Zhang & Ni 2010). We will consider oxygen fugacities around 2 log units below the iron-wustite buffer, as appropriate for the Fe/Mg ratio of our modeled mantles (Frost et al. 2008). This is consistent with lunar basalts and Earth formation models (e.g. Frost et al. 2008, Zhang 2011). For most of our simulated planets there will be enough accreted Fe to react with the water, so our model will produce H2-dominated (low-molecular-weight, low H2O/H2 ratio) atmospheres unless fractionating escape is severe.

Giant impacts can remove the atmosphere. Early analytic studies disagreed on the fraction of the volatile envelope that is lost during giant impacts (Ahrens 1993, Genda and Abe 2003). We use semi-analytic approximations for giant impact stripping fit to CTH simulations (Stewart et al. 2014; see also Schlichting 2015). We use the scaling (Stewart et al.

2014), where b is the impact parameter and (Leinhardt & Stewart 2012). We


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use the log-linear fit indicated by Stewart et al. 2014, with complete atmospheric loss for Qs = 2 × 107 J/kg and complete ocean loss for Qs = 108 J/kg. Planetesimal accretion (and thus outgassing) can continue after the last giant impact.

2.4.5. Confronting the outgassing hypothesis with Kepler (+ K2 + TESS) data. We hypothesize that outgassing can make a major (>50%) contribution to most (>50%) of Kepler’s 1.6 R⊕< R ≲ 2.5 R⊕ planets around FGK stars. To maximize the impact of our research, we will compare the results of the proposed models to astronomical observations, principally exoplanets with measurements of both radii and masses.

Previous work has shown that most planets larger than 1.6 R⊕ are not predominantly rocky, i.e., gas must be responsible for a significant fraction of the planets' radii (Rogers 2015). However, that does not imply that most planets smaller 1.6 R⊕ are predominantly rocky. Other recent research has suggested that planets with radii 1-4 R⊕ typically have 0.5-1% of their mass in H and He (Wolfgang & Lopez 2014). This is within the range allowed by the stoichiometric limit, but may be excluded by the proposed models. Given the inevitable limitations of modeling geophysical processes on unresolved planets, we will focus on comparing observations to our predictions for the largest possible and largest plausible planetary radii (for a given orbital period and hence incident stellar flux) that could have arisen from outgassing.

Planet sizes (relative to their host star) are measured precisely by transit searches by NASA's Kepler, K2 missions, and soon TESS. For some planets around relatively bright stars with favorable temperatures and rotation velocities, the Doppler technique has measured the masses of a few low-mass, transiting planets, mostly at short orbital periods. In order to determine how planets form and how their atmospheres evolve due to external factors (e.g., photoevaporation, tidal or magnetic heating), it is important that we also compare the results of our models to planets at somewhat larger orbital periods (e.g., > 10-15 d). Measuring the masses of such planets will soon be possible thanks to a new generation of Doppler instruments coming online - e.g., Habitable Zone Planet Finder for HET, the Stable High-Resolution Echelle for Keck (SHREK), and the Extreme Precision Spectrograph for Exoplanet Studies (EXPRES). In the mean time, we will be able to compare to mass measurements from transit timing variations (TTVs). TTVs are particularly powerful for measuring the masses of small planets in systems with multiple transiting planets that are tightly packed and/or near low-order mean motion resonances. There are over 100 pairs of strongly interacting planets and 260 planets with detected transit timing variations (Holczer et al. 2015). Co-I Ford is the PI for a recent Exoplanets Research grant to characterize the mass-radius relationship of small planets using transit timing variations using Bayesian Hierarchical modeling to account for measurement uncertainties and detection biases (e.g., Fig. 3 of Wolfgang et al. 2015). That project will derive posterior probability distributions for planet masses and densities for at least several dozen small planets (during Year 1 of the proposed work).

As part of this proposal, we would extend those models to calculate the posterior probability that each small planet's atmosphere could be dominated by H from outgassing. We could find that many planets must have accreted significant gas from the nebula. Alternatively, we could find that a substantial population of planets could have atmospheres dominated by outgassing. Either outcome is intrinsically interesting, as it provides insight into the potential origin and likely composition of the atmospheres of small planets. We will explore how the results vary with the size and orbital period of the planets, as well as the photospheric abundances of their host stars


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(and presumably the disks from which they formed). Of particular, interest we can perform more precise comparisons of the relative properties of multiple planets orbiting a common host star, as this minimizes the effect of uncertainties in stellar properties and properties of the disk from which the planets formed (e.g., Carter et al. 2012, Ciardi et al. 2013).

2.4.6 Assumptions, caveats, and refinements. It is worth emphasizing the limitations and assumptions of these methods. First, this is a data-driven proposal whose primary goal is improved understanding of the overall trends in the volatile-envelope mass among Kepler planets. As with other planet formation models more generally, global constraints on the absolute rates and locations of volatile delivery will require integrating data from mature planets (this proposal) with data from planet-forming disks. Second, we do not track nebular accretion of hydrogen, which must obviously become important at some mass (gas giants are not as common as 1.6 R⊕<R< 2.5 R⊕ objects, but neither are they rare). This means we omit feedbacks between nebular accretion and outgassing that could speed outgassing (for example, a more-extended envelope due to nebular accretion will collect more volatile-rich objects, potentially accelerating thermolysis). We accept this simplification in order to focus our investigation on the less-studied outgassing-only endmember. Third, there are exceptions to the density trends we describe in §2.3 – such as dense-for-its-size Kepler-10c and unusually low-density-for-its-proximity to the star Kepler-11b – and these can be investigated using our results (e.g., Howe & Burrows 2015). Fourth, we assume that oxidation of iron in the nebula (at grain scale, prior to assembly into planetesimals) does not vary strongly from solar conditions (Brearley et al. 2006, see also Ciesla et al. 2003). This means that the model is inapplicable to systems with very different cosmochemistries. Finally, our Fe-Mg-Si-O-H chemistry is simple: in particular, we neglect C, which dominates the volatile envelopes of Venus and Mars. This is acceptable because CO2 envelopes cannot explain the observations. It is easy to refine our model to include more complex chemistry (and varying Fe:Si) because CHIM-XPT is computationally inexpensive and the thermodynamic databases include C (and Al, Ca, …). Our goal is to build a model that treats the bottleneck steps for outgassing in a simple, physically self-consistent way. The build-up and loss of volatile envelopes on Gyr-old small-radius exoplanets has a long and presumably complex history; we propose to test simple hypotheses about that history, with an eye to enabling richer hypotheses and tests in future.

2.5 Perceived Impact of the Proposed Work. Scientific priorities. Our work addresses a key question defined by a recent NASA astrophysics roadmap (NASA Enduring Visions 2013): “What are exoplanets like?” Water-rock interaction is a plausible candidate for the puzzlingly low-densities of low-radius planets discovered by Kepler. Solving this puzzle has implications for the habitability of exoplanets further out – because hydrogen blankets extend the outer limit of the habitable zone (Pierrehumbert & Gaidos (2011), and affect climate stability (Abbot et al. 2011). Maximum hydrogen production leads to coreless exoplanets (Elkins-Tanton & Seager 2008), implying a weak or absent magnetic field (Stevenson 2003, Khurana et al. 2011) and altering magmatic activity (BVSP 1981, Kite et al. 2009). Although the Kepler planets with current density constraints are too hot for life, the processes being investigated (migration, scattering, volatile delivery, and atmospheric growth) occur at radii than span the habitable zone (Dauphas & Morbidelli 2014). Beyond the reach of photoevaporation (period >15 day), trends in volatile envelope-generating processes with increasing orbital separation may provide insights into volatile content of planets in the habitable zone.


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Missions. TESS and the exoplanet component of JWST are missions that target Super-Earths (Ricker et al. 2014), and these worlds appear to be volatile-shrouded on the basis of their densities (Rogers 2015). Because of the enormous difficulty of detecting gases directly (Seager 2014), our approach to understanding the cause of the mean and scatter of the volatile-mass-fraction distribution is relevant to understanding density constraints on TESS planets. Planets with H2-rich envelopes are of particular interest for transit spectroscopy because molecular features are more easily resolved in large-scale-height atmospheres (Benneke & Seager 2012, de Wit & Seager 2013); new telescopes will push these detections to smaller-radius planets.

2.6 Relevance of the Proposed Work. Our proposed work will advance knowledge about the effect of water-rock interaction on the volatile envelopes required to explain observations of Kepler’s planets. Therefore, our proposal is within the scope of the Exoplanets call, specifically “[I]nterpret observations of exoplanetary systems,” “[U]nderstand the chemical and physical processes of exoplanets (including their […] atmospheres),” and “[I]mprove understanding of the origins of exoplanetary systems.”

2.7 Work Plan. Activities/milestones Deliverables

Yea

r 1

• Carry out idealized degassing calculations, forced by idealized (e-folding) migration history without giant impacts.

• Postdoc sets up geochemical code under Kite’s supervision. • Begin analysis of N-body simulations. • Start analysis of Kepler dataset.

✓ LPSC presentation on: Geochemical and geophysical model with analytic accretion. ✓ Short ApJL-length manuscript on: Idealized degassing calculations – what is the efficiency relative to the stoichiometric limit?

Yea

r 2

• Complete idealized water-rock interaction calculations, and carry out realistic water-delivery simulations runs.

• Complete analysis of Kepler dataset, and compare with “optimal” water-delivery results.

• Complete analysis of N-body simulations and derive spread of predictions for giant-impact atmospheric erosion.

• Incorporate ensemble of “realistic” degassing calculations into predictions of contribution of water-rock interaction to observations for Kepler planets.

✓ LPSC presentation on: Hydrogen production and retention under “optimal” water-delivery scenarios. ✓ Detailed manuscript on: Model predictions for “optimal” and “realistic” water-delivery scenarios, including evaluation of model relative to two key planetary systems (Kepler-20 and Kepler-106).

Yea

r 3

• Complete parameter sweeps using the integrated dynamics-outgassing model. Analyze the results and implications the origin of volatile envelopes on Kepler planets.

• Carry out sensitivity tests and refine the model in order to quantify the possible effects of non-standard rock major-element compositions.

✓ LPSC presentation on: Hydrogen production and retention under “realistic” water-delivery scenarios. ✓ Detailed manuscript constraining contribution of outgassing to volatile envelopes on Kepler, K2, and TESS planets.

2.8 Personnel and Qualifications. (For FTE information, see §6, Budget Justification). PI Edwin Kite is an Assistant Professor of Planetary Geoscience at the University of Chicago (from 1 Jan 2015). As PI, he will participate to some degree in all aspects of the proposed work and oversee its implementation. He will develop and test the physical models. Under Kite’s supervision, UChicago Postdoctoral Researcher Mohit Melwani Daswani will lead the implementation of the geochemical models, and will be trained by Kite in the integration of the geochemical models with the planet-resurfacing code. Co-I Eric Ford is a Professor of Exoplanets at Penn State; he will assist with the setup and analysis of the N-body calculations, help with the analysis and interpretation of the output from the combined outgassing-dynamical models, and lead the comparison of simulation results to the properties of planets discovered by NASA missions. All personnel will participate in interpretation of results.


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Schaefer, Laura; Fegley, Bruce (2007), Outgassing of ordinary chondritic material and some of its implications for the chemistry of asteroids, planets, and satellites, Icarus, Volume 186, Issue 2, p. 462-483. Schaefer, Laura; Fegley, Bruce (2009), Chemistry of atmospheres formed during accretion of the Earth and other terrestrial planets, Icarus, Volume 208, Issue 1, p. 438-448. Schlaufman, Kevin C. (2015), A Continuum of Planet Formation between 1 and 4 Earth Radii, The Astrophysical Journal, Volume 790, Issue 2, article id. 91, 11 pp. (2014). Schlichting, Hilke E. (2015), Formation of Close in Super-Earths and Mini-Neptunes: Required Disk Masses and their Implications, The Astrophysical Journal Letters, Volume 795, Issue 1, article id. L15, 5 pp. (2014). Schubert, Gerald; Turcotte, Donald L.; Olson, Peter (2001), Mantle Convection in the Earth and Planets, Mantle Convection in the Earth and Planets, by Gerald Schubert and Donald L. Turcotte and Peter Olson, pp. 956. ISBN 052135367X. Cambridge, UK: Cambridge University Press, September 2001. Seager, S. (2014), The future of spectroscopic life detection on exoplanets, Proceedings of the National Academy of Sciences, vol. 111, issue 35, pp. 12634-12640. Shackelford JF, Studt PL, Fulrath RM (1972) Solubility of gases in glass. II. He, Ne, and H2 in fused silica. J Appl Phys 43:1619-1626. Silburt, Ari; Gaidos, Eric; Wu, Yanqin (2015), A Statistical Reconstruction of the Planet Population around Kepler Solar-type Stars, The Astrophysical Journal, Volume 799, Issue 2, article id. 180, 12 pp. Sleep, N. H.; Meibom, A.; Fridriksson, Th.; Coleman, R. G.; Bird, D. K. (2004), H2-rich fluids from serpentinization: Geochemical and biotic implications, Proceedings of the National Academy of Science, vol. 101, Issue 35, p.12818-12823. Solomon, S. C. (1979), Formation, history and energetics of cores in the terrestrial planets, Lunar and Planetary Institute and NASA, Workshop on Solid Convection in the Terrestrial Planets, Moffett Field, Calif., Dec. 12, 13, 1977. Physics of the Earth and Planetary Interiors, vol. 19, p. 168-182. Stevenson, David J. (2003), Planetary magnetic fields, Earth and Planetary Science Letters, Volume 208, Issue 1-2, p. 1-11. Toksoz, M. N.; Hsui, A. T.; Johnston, D. H. (1978), Thermal evolutions of the terrestrial planets Moon and the Planets, vol. 18, p. 281-320. Vance, Steve; Harnmeijer, Jelte; Kimura, Jun; Hussmann, Hauke; deMartin, Brian; Brown, J. Michael (2007), Hydrothermal Systems in Small Ocean Planets, Astrobiology, Volume 7, Issue 6, pp. 987-1005. Young E. D. (2001) The hydrology of carbonaceous chondrite parent bodies and the evolution of planet progenitors. Phil. Trans. R. Soc. London, A359, 2095–2109. Schlichting, Hilke E.; Sari, Re'em; Yalinewich, Almog (2014), Atmospheric mass loss during planet formation: The importance of planetesimal impacts, Icarus, Volume 247, p. 81-94.


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Zhang , Y., & Huaiwei Ni (2010), Diffusion of H, C, and O Components in Silicate Melts, Reviews in Mineralogy & Geochemistry Vol. 72 pp. 171-225, 2010. DOI: 10.2138/rmg.2010.72.5 Zolensky, M.E., W.L. Bourcier, J.L. Gooding (1989), Aqueous alteration on the hydrous asteroids: results of EQ3/6 computer simulations, Icarus, 78, pp. 411–425 Zolensky, M.E., Alexander N. Krot and Gretchen K. Benedix (2008), Record of low-temperature alteration in asteroids, in Oxygen in the solar system: Reviews in Mineralogy and Geochemistry, volume 68, 429-462.


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4. Biographical Sketches Edwin S. Kite (Principal Investigator). Professional preparation: B.A. Cambridge University (Natural Sciences Tripos – Geological Sciences), June 2007 M.Sci. Cambridge University (Natural Sciences Tripos – Geological Sciences), June 2007 Ph.D. University of California Berkeley (Earth and Planetary Science), December 2011 Appointments: Assistant Professor, Department of the Geophysical Sciences, University of Chicago January 2015 – Harry Hess Fellow (Astrophysical Sciences and Geosciences Departments), Princeton University January 2014 – December 2014 O.K. Earl Fellow, Geological and Planetary Sciences, California Institute of Technology. January 2012 – January 2014. Papers

Borlina, C., Ehlmann, B.E., & Kite, E.S., “Modeling the thermal and physical evolution of Mount Sharp’s sedimentary rocks, Gale Crater, Mars: implications for diagenetic minerals on the MSL Curiosity rover traverse,” moderate revisions requested by J. Geophys. Res. - Planets.

Kite, E.S., Howard, A., & Lucas, A., 2015, “Resolving the era of river-forming climates on Mars using stratigraphic logs of river-deposit dimensions,” Earth Planet. Sci. Lett., 420, 55-65.

Kite, E.S., Howard, A., Lucas, A., Armstrong, J.C., Aharonson, O., & Lamb, M.P., 2015, “Stratigraphy of Aeolis Dorsa, Mars: stratigraphic context of the great river deposits,” Icarus, 253, 223-242.

Kite, E.S., Williams, J.-P., Lucas, A., & Aharonson, O., 2014, “Paleopressure of Mars' atmosphere from small ancient craters,” Nature Geoscience, 7, 335-339.

Kite, E.S., Lewis, K.W., Lamb, M.P., Newman, C.E., & Richardson, M.I., 2013, “Growth and form of the mound in Gale Crater, Mars: Slope-wind enhanced erosion and transport,” Geology, 41, 543-546. (“Highlight of the AGU Meeting” at Science).

Kite, E.S., Halevy, I., Kahre, M.A., Manga, M., & Wolff, M., 2013, “Seasonal melting and the formation of sedimentary rocks on Mars, with predictions for the Gale crater mound" Icarus, 223, 181-210.

Kite, E.S., Lucas, A., & C.I. Fassett, 2013, “Pacing Early Mars river activity: Embedded craters in the Aeolis Dorsa region imply river activity spanned ≳(1-20) Myr,” Icarus, 225, 850-855. Rappaport, S., Levine, A., Chiang, E., El Mellah, I., Jenkin, J., Kalomeni, B., Kite, E.S., Kotson, M., Nelson, L., Rousseau-Nepton, & Tran, K., 2012, “Possible disintegrating short-period Super-Mercury orbiting KIC 12557548," Astrophys. J., 752, 1. Kite, E.S., Gaidos, E. & M. Manga, 2011. “Climate instability on tidally locked exoplanets,” Astrophys. J., 743, 41, 12 pp.


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Kite, E.S., Manga, M., & Gaidos, E., 2009. “Geodynamics and rate of volcanism on massive Earth-like planets," Astrophys. J., 700, 1732-1749.

Kite, E.S., Michaels, T.I., Rafkin, S.C.R., Manga, M., & W.E. Dietrich, 2011. “Localized precipitation and runoff on Mars,” J. Geophys. Res. Planets, 116, E07002, 20 pp. doi:10.1029/2010JE003783.

Kite, E.S., Rafkin, S.C.R., Michaels, T.I., Dietrich, W.E., & Manga, M., 2011. “Chaos terrain, storms, and past climate on Mars,” J. Geophys. Res. Planets, 116, E10002, 26 pp., doi:10.1029/2010JE003792 (“Research Highlight” at Nature Geoscience). Manga, M., Patel, A., Dufek, J., & Kite, E.S., 2012. “Wet surface and dense atmosphere on early Mars inferred from the bomb sag at Home Plate, Mars,” Geophys. Res. Lett., doi:10.1029/2011GL050192.

Mangold, N., Kite, E.S., Kleinhans, M., Newsom, H.E., Ansan, V., Hauber, E., Kraal, E., Quantin-Nataf, C. & K. Tanaka, “The origin and timing of fluvial activity at Eberswalde Crater, Mars,” Icarus, 220, 530-551.

Šrámek, O., McDonough, W., Kite, E.S., Lekic, V., Zhong, S.T., & Dye, W.F., Geophysical and geochemical constraints on geoneutrino fluxes from Earth's mantle, Earth Planet. Sci. Lett., doi:10.1016/j.epsl.2012.11.001. (“Research Highlight” at Nature).

Chiang, E., Kite, E., Kalas, P., Graham, J. R., & Clampin, M., 2009. “Fomalhaut's Debris Disk and Planet: Inferring the Mass and Orbit of Fomalhaut b Using Disk Morphology," Astrophys. J., 693, 734-749.

Kite, E.S., Matsuyama, I., Manga, M., Perron, J.T., & Mitrovica, J.X., 2009, “True polar wander driven by late-stage volcanism and the distribution of paleopolar deposits on Mars," Earth Planet. Sci. Lett, 280, 254-267.

Kite, E.S., & R.C.A. Hindmarsh, 2007. “Did ice streams shape the largest channels on Mars?”, Geophys. Res. Lett., 34, L19202. Kalas, P., Graham, J. R., Chiang, E., Fitzgerald, M. P., Clampin, M., Kite, E. S., Stapelfeldt, K., Marois, C., & Krist, J., 2008. “Optical Images of a Planet 25 Light Years from Earth," Science, 322, 1345-1348 (“Breakthrough of the Year #2” at Science; AAAS Newcomb Cleveland Prize).

Field geology experience: Central India (Proterozoic). Greece, SE Spain, England, Scotland, California, Hawaii (fieldwork, mapping courses). NW Spain (independent mapping project, 6 weeks). Utah (GSI for Professor W. Alvarez). Telescope experience: Hubble Space Telescope: Co-I on GO/DD Program 11818 (PI: Paul Kalas). Spitzer Space Telescope: Warm IRAC phase curves of exoplanet HAT-P-7b (PI: Heather Knutson).

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Biographical Sketch — Eric B. Ford

Research Interests & Expertise Exoplanet Detection & Characterization, Planetary Dynamics & Formation, Astrostatistics, Astrobiology

Education Princeton University, Ph.D. Astrophysical Sciences (2003) Massachusetts Institute of Technology, S.B. Physics (1999), S.B. Mathematics (1999)

Appointments Pennsylvania State University

Professor, Astronomy & Astrophysics and Institute for CyberScience (2013–); Associate Director, Center for Astrostatistics (2013–); Deputy Director, Center for Exoplanets & Habitable Worlds (2013–)

University of Florida Assistant Professor, Astronomy (2007–2010); Associate Professor, Astronomy (2010–2013); Affiliate Associate Professor, Physics (2011–2013)

Harvard-Smithsonian Center for Astrophysics (CfA) NASA Hubble Postdoctoral Fellow, CfA Institute for Theory & Computation (2006–2007)

University of California, Berkeley

Miller Fellow, Astronomy Department & Miller Institute for Basic Research in Science (2003–2006)

Five Selected Recent Publications: (of 126 refereed papers, incl. 24 as first author, 22 with advisees as first author, 20 in Science or Nature; 12,875 citations. h-index: 59; current/recent advisees underlined)

Jontof-Hutter, D., J.F.Rowe, J.J. Lissauer, D.C. Fabrycky, E.B. Ford (2015) accepted to Nature. “Mass of the Mars-sized Exoplanet Kepler-138 b from Transit Timing”

Chatterjee, S. & E.B. Ford (2015) Astrophysical Journal 803, 33. “Planetesimal Interactions Can Explain the Mysterious Period Ratios of Small Near-Resonant Planets”

Boley, A.C., M. Morris & E.B. Ford (2014) Astrophysical Journal, 792, 27. “Overcoming the Meter Barrier and the Formation of Systems with Tightly Packed Inner Planets”

Ford, E.B. (2014) PNAS, 111, 12616. “Architectures of planetary systems and implications for their formation”

Nelson, B., Ford, E.B., Wright, J.T., D.A. Fischer, K. von Braun, A.W. Howard, M.J. Payne, S. Dindar (2014) Monthly Notices of Royal Astronomical Society, 441, 442. “55 Cnc planetary system: fully self-consistent N-body constraints & a dynamical analysis”

Outcomes of Recent, Most Closely Related Research Efforts Ford was PI for one of 9 Participating Scientists on the science team for Kepler, NASA's first

mission to search for Earth-sized planets, contributing to 55 refereed papers and the science team via the Kepler Transit Timing & Multi-body Working Group (2009 – present), Short Cadence Review Team (2012 – 2013), Kepler Science Council (2012), Exoplanet Council (2012 – 2014), and Science Operations Center Data Validation Peer Review Panel (2008).

Ford was PI for a NASA Origins award “Self-Consistent Dynamical Modeling of Multiple Planet Systems,” contributing to 21 refereed publications, including led by 10 Ford and/or his advisees (not including one recently submitted to ApJ and on arXiv). The team analyzed nearly all of the strongly interacting multiple planet systems with publicly available radial velocity measurements, including HD 12661 HIP 14810, HD 11964, HD 38529, HD 108874, HD 168443, HD 190360, HD 108874, HD 168443, 24 Sextans, HD 200964, HD 37124, HD 37605, MARVELS-1b, HD 82943, Rho Cnc and GJ 876.

xrp 2015 origin of the volatile envelopes of small...

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