Ruth Murray-Clay - University of Chicago Murray-Clay Harvard-Smithsonian Center for Astrophysics Image Credit: NASA. Graduate Students Undergraduates Piso Tripathi ... Ram: τ ∝

Origins of Gas Giant Planets

Ruth Murray-ClayHarvard-Smithsonian Center for Astrophysics

Image Credit: NASA

Graduate Students

Undergraduates

Piso Tripathi

Dawson

Wolff Lau Alpert Mukherjee Wolansky Jackson

Marois et al. 2010

Neptune’s orbital distance

HR 8799: A testbed for planet formation theories

NASA, ESA, M. Robberto, HST Orion Treasury Project, L. Ricci

Alves, Lada, & Lada 2001



gas + dust + ice

gas + dust

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*

Data: Zuckerman & Song 2009; exoplanet.eu

Kratter, Murray-Clay, & Youdin, ApJ (2010)

~planets

~brown dwarfs

~stars

Test case HR 8799: Brown Dwarfs or Planets?

Jupiter

Saturn

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*



~planets

~brown dwarfs

~stars


Jupiter

SaturnPlanetary companions

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*



~planets

~brown dwarfs

~stars

Test case HR 8799: Brown Dwarfs or Planets?Brown dwarf companions

Jupiter


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*


Kratter, Murray-Clay, & Youdin, ApJ (2010) Larger than most

protoplanetary disks

~planets

~brown dwarfs

~stars

Test case HR 8799: Brown Dwarfs or Planets?Brown dwarf companions

Jupiter


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

Planetary companions

Brown dwarf companions

HR 8799Jupiter

Saturn


Three ways to form companions

turbulent fragmentation

core accretion

gravitational instability

Three ways to form companions

turbulent fragmentation

core accretion

gravitational instability

HR 8799

Today’s Goal:A framework for identifying giant planet

formation processes.

Marois et al. 2010

using HR 8799 as a test case

Outline:Origin Stories

~70 AU

More complicated hybrid scenariose.g., form a core close to the star, scatter outward, then accrete gas

Formation closer to star + scattering outward

~ in situ formation by gravitational instability

~ in situ formation by core accretion


~70 AU





Gravitational instability

L

most unstable scale:

self-gravity overcomes pressure and Keplerian shear

L ⇠ H

Mfrag ⇠ ⌃(2⇡H)2 ⇠ 4⇡

✓H

a

◆3

M⇤

Gravitational Instability?Planets cannot grow after fragmentation and they must migrate in

Kratter, Murray-Clay, & Youdin, ApJ 2010

e10

Mas

s (M

Jup)

100

20 1401

Radius (AU)60 100

minimumfragment mass

minimumfragment distance

Gravitational Instability?Planets cannot grow after fragmentation and they must migrate in

Kratter, Murray-Clay, & Youdin, ApJ 2010

e10

Mas

s (M

Jup)

100

20 1401

Radius (AU)60 100

minimumfragment mass

minimumfragment distance

Stability requires resonance, suggesting migration

Fabrycky & Murray-Clay 2010

Collapse must occur at the end of infallor the fragment will grow into a binary star

M

Mdisk


M

Mdisk


M

Mdisk


M

isolation mass is stellar!

Mdisk

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

HR 8799expect higher

mass companions from disk collapse

Gravitational instability planets can only be failed binary stars


~70 AU





Core accretion: Need to grow a solid core through collisions

v t

v

Volume = Avt

# collisions in time t ~ n Avt ~ (nH) A (t/torb)

number density = n

A

H ~ v torb

Sun

Growth timescale is set by:A: cross-section for collisionsv: relative velocities of colliding bodiesn: density of bodies available to accrete

mass of object trying to grow

orbital time: longer farther from the star

cross-sectionfor collisions

disk surface density in solids

tgrow

=M

M⇠ M

⌃Atorb

mass of object trying to grow

orbital time: longer farther from the star

cross-sectionfor collisions

disk surface density in solids

tgrow

=M

M⇠ M

⌃Atorb

Cross-section regimes:

physical cross-sectiongravitational focusing

capture by gas drag intoplanetary atmosphere ?

Sun

Let’s make gas useful

planetesimal

larger planetesimal or core

Gas Alters the Orbits of Single Planetesimals

planetesimal wants to orbitstar at

radial pressuregradient

but gas orbits more slowly

small

The resulting drag acceleration

large

isvKep

vorb

< vKep

v2orb

r=

GM⇤r2

+1

⇢

dP

dr

FD

m

In the absence of gas, satellites can orbit within the Hill radius

Sun

RHill


Sun

RHill

planetarygravity


Sun

RHill

planetarygravity


Sun

RHill

planetarygravity


Sun

RHill

tidalacceleration

planetarygravity


Sun

RHill

tidalacceleration

planetarygravity

RHill

RWISH

RHill

No gas: Gas:

core

planetesimal

Wind Shearing (WISH)

Perets & Murray-Clay (2011)

A core + planetesimal in gas can:shear apart

or inspiral and merge

orbit around starmutualorbit




FD

m

FD

m

10−4 10−2 100 102 104 106 108

rs (cm)

10−1

100

101

102

103

104

105

R sta

b/rb

rs ≤ rb

rb/rb

RH/rb

RWS/rb

Epstein

Stokes

1 < Re < 800Ram

1 AU rb = 10kmRam Pressure Drag

binarystability

10−4 10−2 100 102 104 106 10810−1

100

101

102

103

104

105

Wind-Shearing Limits Orbital Stability

WISH radiuslarge body radius

small body radius (cm)

Hill radius

core radius

10km= core radius





FD

m

FD

m

100

102

104

106

108

1010

τm

erge

(yr)

Porb

1 AU

Stokes

Ram: τmerge ∝ rs

100

102

104

106

108

1010

τm

erge

(yr)

Porb

5 AU

Epstein

Stokes

Ram

10−4 10−2 100 102 104 106 108

rs (cm)

100

102

104

106

108

1010

τm

erge

(yr)

Porb

40 AU

Epstein: τmerge ∝ rs

Stokes: τmerge ∝ rs

2

100

102

104

106

108

1010

τm

erge

(yr)

Porb

1 AU

Stokes

Ram: τmerge ∝ rs

100

102

104

106

108

1010

τm

erge

(yr)

Porb

5 AU

Epstein

Stokes

Ram

10−4 10−2 100 102 104 106 108

rs (cm)

100

102

104

106

108

1010

τm

erge

(yr)

Porb

40 AU



2

100

102

104

106

108

1010

τm

erge

(yr)

Porb

1 AU

Stokes

Ram: τmerge ∝ rs

100

102

104

106

108

1010

τm

erge

(yr)

Porb

5 AU

Epstein

Stokes

Ram

10−4 10−2 100 102 104 106 108

rs (cm)

100

102

104

106

108

1010

τm

erge

(yr)

Porb

40 AU



2

small body radius (cm)

For Small Planetesimals,Merger Times < Disk Lifetimes

time

to m

erge

(yr

)

typical disk lifetime

~Pluto radius ~km


“Binary Capture”

RHill

RWISH

dissipation due to interaction

with gas

Murray-Clay & Perets (in prep)(see also Ormel & Klahr 2010)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02

y (A

U)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02

y (A

U)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02y

(AU

)

RHill

Ratm

Integration of forces confirms that capture can happen with a cross-section as large as the Hill radius

RHill

RWISH

Ratm Capture by the atmosphere is only possible if

the small particle can decouple

from the exterior gas.

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02y

(AU

)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02

y (A

U)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02

y (A

U)

RHill

Ratm

RWISH

Integrations confirm that well-entrained particles can be swept around the core, preventing accretion

Growth times at 70 AU can be short enough to nucleate an atmosphere

104 105 106

time (yr)

10ï410ï310ï210ï1100101102103

Mco

re (M

Earth

) Mcrit(Rafikov

2006,2010)

50% of MMSN solids in equal mass per log bin from mm

to 10cm

Murray-Clay et al. (in prep.)

Growth times at 70 AU can be short enough to nucleate an atmosphere

104 105 106

time (yr)

10ï410ï310ï210ï1100101102103

Mco

re (M

Earth

) Mcrit(Rafikov

2006,2010)

50% of MMSN solids in equal mass per log bin from mm

to 10cm

Murray-Clay et al. (in prep.)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02

y (A

U)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02

y (A

U)

ï0.02ï0.01 0.00 0.01 0.02x (AU)

ï0.02

ï0.01

0.00

0.01

0.02y

(AU

)

� = 0.01

RHill

Ratm

Turbulence doesn’t prevent the final stage of core growth by capture of planetesimals

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

HR 8799: Brown Dwarfs or Planets?

HR 8799

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars


HR 8799

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars


HR 8799

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

HR 8799

Gemini Planet Imager (PI: Bruce Macintosh)

will have different mass distributionsOverlapping populations

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100

Mp/M

*


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100

Mp/M

*


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100

Mp/M

*


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100

Mp/M

*

Num

ber

Mass

Num

ber

Mass

Excluding radial orbits:

Stellar metallicity and planetary atmospheric composition provide additional discriminants

H2O snowline ~2 AU

CO snowline ~40 AU C and O

frozen outtogether

C rich gasO rich solids

Oberg, Murray-Clay, & Bergin, ApJL

(2011)

O

C


~70 AU





Dawson & Murray-Clay 2013

~ in situ formation?hot Jupiters

highly eccentric planets

Metal-rich stars host more “moved” planets: A signature of planet-planet interactions

“moved” planets{

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

HR 8799

GPI

Jupiter

Saturn


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

HR 8799

GPI

Jupiter

Saturngravitational instability


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

HR 8799

GPI

Jupiter

Saturncore

accretion


100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

core accretion

disk gravitational instability?

HR 8799: signpostJupiter

Saturn

How do giant planets and brown dwarfs form?turbulent fragmentation

Formation and evolution of planetary systems

Infall and disk accretion sets the

initial conditions for

and then evolve dynamically over long

timescales

which then migrate as they interact with other planets,

residual gas and planetesimals

around stars with different masses, at different stages of evolution, and in

different environments.

growth of planetesimals and

planets

all the while developing

atmospheres that depend on this

history.

Tripathi

Wolff, Dawson, & Murray-Clay (2012)

Piso, Youdin, & Murray-Clay (final stages)

Dawson & Murray-Clay (2013)

Lau & Murray-Clay (final stages)

Dawson & Murray-Clay (2012)

Alpert

Mukherjee Dawson, Murray-Clay, & Johnson (submitted)

Dawson et al. (2012)

Dawson, Murray-Clay, & Fabrycky (2011)Graduate StudentsUndergraduates

Wolansky (senior thesis)

Jackson (junior project)

100 102 104

rp (AU)

10ï4

10ï3

10ï2

10ï1

100M

p/M*




~planets

~brown dwarfs

~stars

core accretion

disk gravitational instability?

HR 8799: signpostJupiter

Saturn

How do giant planets and brown dwarfs form?turbulent fragmentation

Ruth Murray-Clay - University of Chicago Murray-Clay Harvard-Smithsonian Center for Astrophysics Image Credit: NASA. Graduate Students Undergraduates Piso Tripathi ... Ram: τ ∝

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