Origin of the terrestrial planet‘s water, crust and atmosphere

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

2. 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?

3. Interior Atmosphere Interaction - Volcanic outgassing(secondaryatmosphere) 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

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

5. Effect of Water on mantleviscosity (thus on strength of convection and coolingefficiency) Newtonian 1021 typical values for E and V 1019 >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

6. Terrestrialplanetshad (some still have) water in theirinterior • Planet withstagnantlidconvectionmay not beableto bring waterintotheinteriorlate in theevolution • Plate-tectonic planet transports volatiles into the interior by subduction, but does PT first need water inside?

7. 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

8. 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.%

9. 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

10. Water supply from comets less likely

11. From where does the water come in the first place? Do we start with dry or ‘wet‘ bodies? • Accretionmodelssuggestearlymixing of water-richplanetesimalsintotheinner solar system (e.g., Walsh et al. 2011, O‘Brien et al. 2014)

12. From where does the water come in the first place? Do we start with dry or ‘wet‘ bodies? • Accretionmodelssuggestearlymixing of water-richplanetesimalsintotheinner solar system (e.g., Walsh et al. 2011, O‘Brien et al. 2014) • Ruand oxygen isotopes of lunar and Earth samplesaswellasanalysis of angrites and eucritessuggestwater in theinner solar systemwithinthefirstfew Ma (e.g., Sarafinet al, 2014, 2017; Greenwood et al. 2018)  Accretion of dry and volatile-richplanetsimals

13. 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

14. What processes determined the amount of volatiles in the planetary interior? • Accretion of planetary material from planetarynebula: volatile composition of primordial material forminga planet • Early catastrophicwaterloss/ outgassing due to • Oxidation • Impact dehydration • Magma oceansolidification • Degassingbysecondaryvolcanism • Recyling of volatiles bytectonics

15. 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.

16. 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

17. 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 (Tyburczy et al. 2001)

18. 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

19. Magma ocean crystallization and outgassing • At latestage of accretion, an earlymagmaoceancouldhavebeenpresent melting temperature Temperature (K) Accretional temp. profile Radius

20. Magma Ocean Crystallization • Freezing of a magmaoceanresults in a fractionatedmantlewhichisunstabletogravitationaloverturn Cooling Bulk Mg # adiabats liquid solidus <25 liquidus Liquid silicate liquid + solid Radius 89 Solidified magma ocean 85 Enrichement in Fe 81 solid Temperature 84 [Elkins-Tanton et al., 2005]

21. Earth: Cumulate mantle density before and after overturn for a 2000 km deep terrestrial magma ocean Tikoo and Elkins-Tanton (2017)

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

24. Magma ocean crystallization and outgassing • Whenfreezing from bottomto top, volatiles (simlartoradioactiveelements) arecontinouslyenriched in themelt • Efficientand rapid degassing: meltreachesbyvigorousconvectionthesurfacewheresaturation of volatiles in themeltislow Lebrun et al., 2013

25. Tikoo and Elkins-Tanton (2017)

26. 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)

27. Tikoo and Elkins-Tanton (2017)

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

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

30. Water content of mantle cumulates from equilibrium partitioning between magma ocean liquids and fractionating solids. Tikoo and Elkins-Tanton (2017)

31. Tikoo and Elkins-Tanton (2017)

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

33. 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

34. Initial water distribution after core formation and magma ocean crystallization • Most interiorwaterdegassed in theearlystages but a ‘critical‘ amountcanremain in theinterior • Volatiles remain in themantle in trappedmelt -- valuesvarybetween 1-10 % (Elkins-Tanton et al. 2008; Hier-Majumder and Hirschmann 2014) • Different scenarios of waterdistribution in theinterior

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

36. Evolution of early atmosphere Nikolaou et al. 2018

37. 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

38. 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

39. Take home message • Accretion of dry and water-richplanetesimals • Strong devolatilitzationof theinteriorcanbeexpectedearly in evolution (oxidation, impactdevolatization and magmaocean outgassing) but thedensertheearlyatmosphere and thefasterthesolidificationthe morewatercanremain in theinterior (ppm level) accretion of dry and wet planetsimals and inefficient oxidation and degassing - part of water remains in the interior

40. 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)