Origin of the terrestrial planet‘s water, crust and atmosphere D. Breuer (DLR, Berlin) Summer School “Basics of Astrobiology“, Vienna

Outline From where and when does the water come in the first place? Processes of volatile loss and input How is water distributed after planet formation? interior – ocean – (early) atmosphere Do we start with dry or ‘wet‘ interiors?

Interior Atmosphere Interaction - Volcanic outgassing (secondary atmosphere) Volatiles are soluted in the melt Melt is transported into the crust due to its buoyancy. Some of the melt will erupt extrusively. Volatiles embedded in the melt will be partly outgassed Atmosphere is enriched in volatiles

Effect of Water on Solidus (thus on volcanic activity) [after Katz et al. 2003] Mantle water content has a large influence on the mantle solidus 100 ppm water reduce the solidus by ~ 30K.

typical values for E and V Effect of Water on mantle viscosity (thus on strength of convection and cooling efficiency) Newtonian typical values for E and V 1019 1021 >1040 Pas E = 300 – 540 kJ/mol V = 2·10-6 – 2·10-5 m3/mol depth Small amount of water (~100 ppm) reduces viscosity by two orders of magnitude dry wet

Terrestrial planets had (some still have) water in their interior Planet with stagnant lid convection may not be able to bring water into the interior late in the evolution Plate-tectonic planet transports volatiles into the interior by subduction, but does PT first need water inside?

Possible accretion scenarios dry accretion and late supply of volatile-rich planetesimals accretion of dry and wet planetsimals but water is efficiently removed from the interior by oxidation, impact and magma ocean degassing accretion of dry and wet planetsimals and inefficient oxidation and degassing - part of water remains in the interior

What is wet and what is dry? Ganymede ~ 50 wt.% ice/water fraction Earth ~ 0.03 wt.% water fraction (1 surface ocean) ~ 0.015 – 0.15 wt.% (0.5 – 5 ocean masses in the interior) Ocean is equal to 270 bar atmosphere Mars present-day dry surface but indication for past water at the surface and ‘wet‘ interior ~ 0.01 wt.%

When inner planets are formed from buildings blocks of their ‘original‘ location they should be dry 100 10 Water content wt. % 1 0.1 0.01

Water supply from comets less likely

From where does the water come in the first place From where does the water come in the first place? Do we start with dry or ‘wet‘ bodies? Accretion models suggest early mixing of water-rich planetesimals into the inner solar system (e.g., Walsh et al. 2011, O‘Brien et al. 2014)

From where does the water come in the first place From where does the water come in the first place? Do we start with dry or ‘wet‘ bodies? Accretion models suggest early mixing of water-rich planetesimals into the inner solar system (e.g., Walsh et al. 2011, O‘Brien et al. 2014) Ru and oxygen isotopes of lunar and Earth samples as well as analysis of angrites and eucrites suggest water in the inner solar system within the first few Ma (e.g., Sarafin et al, 2014, 2017; Greenwood et al. 2018)  Accretion of dry and volatile-rich planetsimals

Possible accretion scenarios dry accretion and late supply of volatile-rich planetesimals accretion of dry and wet planetsimals but water efficiently removed from the interior by oxidation, impact and magma ocean degassing accretion of dry and wet planetsimals and inefficient oxidation and degassing - part of water remains in the interior

What processes determined the amount of volatiles in the planetary interior? Accretion of planetary material from planetary nebula: volatile composition of primordial material forming a planet Early catastrophic water loss/ outgassing due to Oxidation Impact dehydration Magma ocean solidification Degassing by secondary volcanism Recyling of volatiles by tectonics

Volatile loss during accretion (1) Oxidation H2O added to inner planets during their accretion is converted on reaction with metallic Fe to FeO and H2 of which the former remained in the mantle and the latter form early atmosphere and escaped.

Oxidation (before core formation) to be efficient: requires homogeneous accretion, i.e., the dry and the volatile–rich components are delivered almost at the same time and before core formation were completed. at high temperatures and pressures in terrestrial magma oceans, iron moves into a metallic phase preferentially to the oxidized phase, eliminating the possibility of significant oxidation. Inefficient oxidation when large iron blobs sink rapidly forming the core

Volatile loss during accretion (2) Impact (shock-induced) devolatilization As a planet grows, the impact velocities increase owing to the increase in radius until first partial and then, at a larger radius, complete devolatilization occurs. Estimates for complete volatilization when radius of target larger than ~1400 km to 2000 km  Decrease of water content toward the surface Comparison Earth versus Mars, small bodies may devolatilize less volatiles (Tyburczy et al. 2001)

New impact experiments This impactor-derived water (about 30%) resides in two distinct reservoirs: in impact melts and projectile survivors. Impact melt hosts the bulk of the delivered water (Daly & Schultz 2018)  growing terrestrial planets may trap water in their interior as the grow

Magma ocean crystallization and outgassing At late stage of accretion, an early magma ocean could have been present melting temperature Temperature (K) Accretional temp. profile Radius

Magma Ocean Crystallization Freezing of a magma ocean results in a fractionated mantle which is unstable to gravitational overturn <25 84 81 89 85 Bulk Mg # Solidified magma ocean Liquid silicate solid liquid + adiabats solidus liquidus Cooling Temperature Radius Enrichement in Fe [Elkins-Tanton et al., 2005]

Earth: Cumulate mantle density before and after overturn for a 2000 km deep terrestrial magma ocean Cumulate mantle density before and after overturn for a 2000 km deep terrestrial magma ocean with an initial water content of 0.25 wt% at solidus temperatures and a reference pressure of 1 atm. Tikoo and Elkins-Tanton (2017)

Mars: density profile before and after overturn Plesa et al. 2015 Elkins-Tanton et al. 2003

Magma ocean crystallization and outgassing When freezing from bottom to top, volatiles (simlar to radioactive elements) are continously enriched in the melt Efficient and rapid degassing: melt reaches by vigorous convection the surface where saturation of volatiles in the melt is low Lebrun et al., 2013

Tikoo and Elkins-Tanton (2017)

Partitioning of water in melt Blue: Batch Melting Red: Fractional Melting Partition coefficient is D=0.01 for water (Katz et al., 2003) Solubility generally higher than partitioning into melt Melt concentration depends linearly on mantle concentration ~0.1 wt.% for F=10% (100 ppm) Y-Achsen Einheit!!

Tikoo and Elkins-Tanton (2017)

Water distribution in cumulates (after overturn) Elkins-Tanton 2008

Water storage in the mantle Upper mantle Transition zone Lower mantle (Hirschmann, 2006)

Tikoo and Elkins-Tanton (2017) Water content of mantle cumulates from equilibrium partitioning between magma ocean liquids and fractionating solids. Water content of mantle cumulates from equilibrium partitioning between magma ocean liquids and fractionating solids. Interstitial liquids are not included in the model. Model results are shown for a 2000 km deep terrestrial magma ocean with an initial water content of 0.25 wt%. Light grey lines denote cumulate water content before overturn. Black lines depict the pre-overturn water contents of modelled cumulate layers, displayed at the post-overturn depth of the modelled layers. Driven by their high density, these water-rich layers would have dewatered as they sank through the transition zone. Radius ranges with two post-overturn values are regions where cumulates from two initial depths (and thus different compositions and mineralogies) have the same densities, and would settle adjacent to each other at the same radius range on some wavelength. (Online version in colour.) Tikoo and Elkins-Tanton (2017)

Tikoo and Elkins-Tanton (2017)

Initial water distribution after magma ocean crystallization Most interior water degassed in the early stages but a ‘critical‘ amount can remain in the interior

Initial water distribution after magma ocean crystallization Trapped melt: compaction rate versus cooling rate The faster the cooling rate the more melt remains in the interor trapped melt

Initial water distribution after core formation and magma ocean crystallization Most interior water degassed in the early stages but a ‘critical‘ amount can remain in the interior Volatiles remain in the mantle in trapped melt -- values vary between 1-10 % (Elkins-Tanton et al. 2008; Hier-Majumder and Hirschmann 2014) Different scenarios of water distribution in the interior

Formation of dense atmosphere: Efficiency of Magma ocean outgassing Solubility of water depends strongly on pressure Gaillard et al. 2013

Evolution of early atmosphere Nikolaou et al. 2018

Main processes for atmosphere formation Capture and accumulation of gasses from planetary nebula Release of H2 due to oxidation of iron Impact volatilization Catastrophic outgassing due to magma ocean solidification Degassing by secondary volcanism Main processes for atmosphere loss Thermal loss process (EUV radiation) Non-thermal loss processes (sputtering, ion loss, photo dissociation ..) Impacts Carbonate formation Absorption in regolith

Water during accretion and MO solidification Oxidation Efficient oxidation before and during core formation (release of H2 into atmopshere) Inefficient oxidation when large iron blobs sink rapidly forming the core Impact devolatization Decrease of water content toward the surface 30 % of volatiles can be stored in impact melt and breccias Magma ocean soldification Efficient degassing first of CO2 and later of H2O Partioning of volatiles into crystalls and trapped melt may result in a damp mantle

Take home message Accretion of dry and water-rich planetesimals Strong devolatilitzation of the interior can be expected early in evolution (oxidation, impact devolatization and magma ocean outgassing) but the denser the early atmosphere and the faster the solidification the more water can remain in the interior (ppm level) accretion of dry and wet planetsimals and inefficient oxidation and degassing - part of water remains in the interior

Take home message Early atmosphere changes from reducing (during core formation) to more oxidized conditions Assuming global magma ocean and fractional crystallization Inhomogeneous distribution of volatiles Water increases toward surface after solidification Remixing of water during MO overturn Potential storage in transition zone (Earth, Venus) and deep mantle (Mars) or homogenous distribution of volatiles? Recyling of volatiles with PT in the case of Earth but initiation of PT is not clear One-way outgassing of interior for stagnant lid planets (Mars, Mercury) Early curst is iron-rich (apart from the Moon (Mercury?) with its plagioclase crust)