Internal structures and compositions of (giant) exoplanets Tristan - PowerPoint PPT Presentation

Internal structures and compositions of (giant) exoplanets Tristan Guillot (OCA, Nice) Exoplanets in Lund Lund 6-8 May 2015

Linking interior & atmospheric composition Interior Atmosphere If c lo an 209458b

Linking interior & atmospheric composition n n o o c c o o r r e e 15 9 cm n o T T e e q q Radius / 10 c = = 10 o 2 2 r 0 0 e 0 0 0 0 K K Interior Atmosphere H-He, no core T = 1 0 0 0 K e e q r o d c e 100 M ⊕ core 100 M ⊕ core a t 5 l M ⊕ 15 M ⊕ core o i s 0 0 1 If brown c planets stars dwarfs 0 lo 0.1 1.0 10.0 100.0 an Mass / M Jup 209458b Moutou et al. (2013) Madhusudhan et al. (2014)

Outline • Jupiter & Saturn • Theoretical considerations • Giant exoplanets • Perspectives

Jupiter & Saturn

Constraints on atmospheric composition

Interior compositions 165-170 K Molecular H 2 (Y~0.23) 1 bar Helium rain 135-145 K 6300-6800 K 1 bar 2 Mbar Molecular H 2 (Y~0.20?) + Metallic H (Y~0.27) 5850-6100 K Helium rain 2 Mbar Metallic + H (Y~ 0.30?) 8500-10000 K 10 Mbar 15000-21000 K 40 Mbar Ices + Rocks core ? Jupiter Saturn

Interior compositions 30 Forbidden region "Slow" "Slow" 25 1 Mbar 1 Mbar 2 Mbar 2 Mbar 8xSolar(Z elements) M core /M Earth 20 3 Mbar 3 Mbar 15 10 atm + 8xSolar(H 2 O) 4 Mbar 4 Mbar a t 5 m 1 Mbar o s 2 Mbar p h e "Fast" 0 r i c 0 2 4 6 8 10 M Z /M Earth

Interior compositions 30 Forbidden region "Slow" "Slow" 25 1 Mbar 1 Mbar 2 Mbar 2 Mbar 8xSolar(Z elements) M core /M Earth 20 3 Mbar 3 Mbar 15 10 atm + 8xSolar(H 2 O) 4 Mbar 4 Mbar a t 5 m 1 Mbar o s 2 Mbar p h e "Fast" 0 r i c 0 2 4 6 8 10 M Z /M Earth Fortney & Nettlemann (2010) Helled & Guillot (2013)

Interior compositions 30 Forbidden region "Slow" "Slow" 25 1 Mbar 1 Mbar 2 Mbar 2 Mbar 8xSolar(Z elements) M core /M Earth 20 3 Mbar 3 Mbar 15 10 atm + 8xSolar(H 2 O) 4 Mbar 4 Mbar a t 5 m 1 Mbar o s 2 Mbar p h e "Fast" 0 r i c 0 2 4 6 8 10 M Z /M Earth

Interior compositions 30 Forbidden region "Slow" "Slow" 25 1 Mbar 1 Mbar 2 Mbar 2 Mbar 8xSolar(Z elements) M core /M Earth 20 3 Mbar 3 Mbar 15 10 atm + 8xSolar(H 2 O) 4 Mbar 4 Mbar a t 5 m 1 Mbar o s 2 Mbar p h e "Fast" 0 r i c 0 2 4 6 8 10 M Z /M Earth

Interior compositions 30 Forbidden region "Slow" "Slow" 25 1 Mbar 1 Mbar 2 Mbar 2 Mbar 8xSolar(Z elements) M core /M Earth 20 3 Mbar 3 Mbar 15 10 atm + 8xSolar(H 2 O) 4 Mbar 4 Mbar a t 5 m 1 Mbar o s 2 Mbar p h e "Fast" 0 r i c 0 2 4 6 8 10 M Z /M Earth

Interior compositions 30 Forbidden region "Slow" "Slow" 25 1 Mbar 1 Mbar 2 Mbar 2 Mbar 8xSolar(Z elements) M core /M Earth 20 3 Mbar 3 Mbar 15 10 atm + 8xSolar(H 2 O) 4 Mbar 4 Mbar a t 5 m 1 Mbar o s 2 Mbar p h e "Fast" 0 r i c 0 2 4 6 8 10 M Z /M Earth

Theoretical considerations

The low-Z content of accreted gas

The low-Z content of accreted gas In standard core-accretion models , most of the heavy elements are accreted during the core growth phase. (A small fraction is accreted during the envelope collapse phase, when the increased gravitational reach brings a fresh supply of planetesimals). Pollack et al. (1996)

The low-Z content of accreted gas In standard core-accretion models , most of the heavy elements are accreted during the core growth phase. (A small fraction is accreted during the envelope collapse phase, when the increased gravitational reach brings a fresh supply of planetesimals). Pollack et al. (1996) With pebble accretion , pebbles are efficiently Log (M protoplanet ) [M Earth ] accreted until the planet reaches the pebble isolation 2 1.0 30 mass (~20 M Earth ). The rest of the accretion then 10 0.1 most of the heavy elements are accreted during the 0 5 r Xfilter=1 [AU] core growth phase. (A small fraction is accreted 3 during the envelope collapse phase, when the 10.0 -2 1 1.0 increased gravitational reach brings a fresh supply of 0.5 planetesimals). (see Lambrechts et al. 2014) -4 0.3 1.0 1 0 . 0.1 -6 -4 -2 0 2 4 Log(Dust size) [cm] Guillot et al. (2014)

The low-Z content of accreted gas In standard core-accretion models , most of the heavy elements are accreted during the core growth phase. (A small fraction is accreted during the envelope collapse phase, when the increased gravitational reach brings a fresh supply of planetesimals). Pollack et al. (1996) With pebble accretion , pebbles are efficiently Log (M protoplanet ) [M Earth ] accreted until the planet reaches the pebble isolation 2 1.0 30 mass (~20 M Earth ). The rest of the accretion then 10 0.1 most of the heavy elements are accreted during the 0 5 r Xfilter=1 [AU] core growth phase. (A small fraction is accreted 3 during the envelope collapse phase, when the 10.0 -2 1 1.0 increased gravitational reach brings a fresh supply of 0.5 planetesimals). (see Lambrechts et al. 2014) -4 0.3 1.0 1 0 . 0.1 -6 Once the planet is formed, the efficiency of -4 -2 0 2 4 Log(Dust size) [cm] planetesimal capture drops (e.g., Guillot & Gladman Guillot et al. (2014) 2000, Matter et al. 2009)

Enrichment of the envelopes • Core accretion: planetesimals are delivered onto the central core. • Core accretion: planetesimals cannot reach the core intact. (Podolak et al. 1988; Pollack et al. 1996) • Envelope capture: accretion efficiency drops (Guillot & Gladman 2000): core erosion? • Present: enriched atmospheres.

Core erosion A 5-20 MEarth core is expected for Jupiter, Saturn, Uranus and Neptune from formation models (see e.g. Mizuno 1980, Ikoma et al. 2001) Is the energy required to erode a primodial core available?

Core erosion A 5-20 MEarth core is expected for Jupiter, Saturn, Uranus and Neptune from formation models (see e.g. Mizuno 1980, Ikoma et al. 2001) Is the energy required to erode a primodial core available? Energy required to mix the core upward:

Core erosion A 5-20 MEarth core is expected for Jupiter, Saturn, Uranus and Neptune from formation models (see e.g. Mizuno 1980, Ikoma et al. 2001) Is the energy required to erode a primodial core available? Energy required to mix the core upward: Maximal core mass flux given intrinsic luminosity L 1 (t): ϖ ≈ 3/10 χ ≈ 0.1: assumes that 10% of the energy in the first convective cell is used to mix chemical elements

Core erosion A 5-20 MEarth core is expected for Jupiter, Saturn, Uranus and Neptune from formation models (see e.g. Mizuno 1980, Ikoma et al. 2001) Is the energy required to erode a primodial core available? Energy required to mix the core upward: Jupiter Maximal core mass flux given intrinsic luminosity L 1 (t): Saturn ϖ ≈ 3/10 χ ≈ 0.1: assumes that 10% of the energy in the first Guillot, Stevenson, Hubbard & Saumon (2004) convective cell is used to mix chemical elements

Core erosion Are elements in the core miscible with metallic hydrogen?

Core erosion Are elements in the core miscible with metallic hydrogen? Water Wilson & Militzer (2011)

Core erosion Are elements in the core miscible with metallic hydrogen? Water Silicates Wilson & Militzer (2011) Wilson & Militzer (2012)

Core erosion Are elements in the core miscible with metallic hydrogen? Water Silicates Wilson & Militzer (2011) Wilson & Militzer (2012) Core erosion is possible because the elements involved are miscible in metallic hydrogen

The noble gases (+N 2 ) Thermodynamic path of the Solar nebula between 5 and 20 AU Delivered with ices as clathrates PH 3 -5.67H 2 O Xe-5.75H 2 O Kr Ar

The noble gases (+N 2 ) Thermodynamic path of the Solar nebula between 5 and 20 AU Delivered with ices as clathrates PH 3 -5.67H 2 O Gautier et al. (2001), Alibert et al. Xe-5.75H 2 O (2005), Mousis et al. (2009) Kr Ar Guillot & Hueso (2006)

The noble gases (+N 2 ) Thermodynamic path of the Solar nebula between 5 and 20 AU Delivered with ices as clathrates PH 3 -5.67H 2 O Gautier et al. (2001), Alibert et al. Xe-5.75H 2 O (2005), Mousis et al. (2009) Kr or Ar Guillot & Hueso (2006)

The noble gases (+N 2 ) Thermodynamic path of the Solar nebula between 5 and 20 AU Delivered with ices as clathrates PH 3 -5.67H 2 O Gautier et al. (2001), Alibert et al. Xe-5.75H 2 O (2005), Mousis et al. (2009) Kr or Ar H-He photoevaporation Disk enriched by H-He 3b K 0 photoevaporation 0 6 ~ 0 5 ~ T H-He photoevaporation 1 3a 0 K 0 , 0 1 0 ~ Guillot & Hueso (2006) T T~10-30K T~100K 2 4 see also Throop & Bally (2010) 1 Low-temperature grains capture gases and settle to the disk mid-plane. Grains migrate in. Some volatiles may be released, but they do not reach the higher altitudes of the disk 2 due to the negative temperature gradient there. The upper atmosphere of the disk evaporates due to radiation from the parent star (3a) and from external radiations (3b). 3 This upper atmosphere contains moslty hydrogen and helium. 4 Giant protoplanets gradually capture a disk gas which is enriched in non-hydrogen-helium species.

Interior vs. Atmosphere