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Planet Formation From Dust to PlanetesimalsCoagulation of dust grains forms millimeter-and centimeter-sized objects

(Dominik and Tielens, 1997)

Stop

p ing

tim

e

Physical Characteristics

Drag forcesStrong -> Week

Chemical bindingStrong -> Week

Surface gravityWeek -> Strong

MAGIC

Collisional coalescence of cm-sized particlesresults in the formation of larger objects andeventually planetesimals (km-sized).

Formation of Terrestrial Planets

Runaway Growth

Gravitational interactioncauses collisions amongplanetesimals and results in the formation of Mercury-to Mars-sized objects (planetary embryo)

Ida and Makino (1993)

Kokubo and Ida (1995, 1996, 1998)

Planetary embryos are formed in ~10,000 y, separated by a few mutual Hill radii.

Accretion of embryos is a local process.

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Final StageGiant impacts among high velocity embryos that result

in terrestrial planets in ~100 million years.

What if the embryos existed also in the asteroid belt?

What if the embryos existed also in the asteroid belt? Water & Earth

- Current location of Earthtoo close to the Sun toretain water

-The icy bodies appearat distances of 4.0 AUand larger

-Earth must have acquiredits water from largerdistances

The variation of relative water content with distance from the Sun implies that water should have been accreted from distant material.

ASTEROIDS (2.5-4.2 AU) OR COMETS (> 30 AU)?

The D/H ratio of Earth’s water rules out a dominant contribution of comets and suggests an asteroidal origin

Numerical integrations also show that comets could have contributed at most 10% of the current water on Earth

Courtesy of F. RobertD/H (x 10-6)

Halley 260-350Hyakutake 280-300Hale Bopp 250-410

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According to asteroid belt sculpting scenario, only 0.1% of the “primitive” asteroids would have been accreted by the Earth. Assuming 1 Earth mass of material and 10% water content this amounts to only 20% of the water currently on Earth. Moreover it arrived “early” in the Earth formation history

WATER FROM ASTEROIDS Water Delivery

• Earth is dry, ~0.05% H2O by mass.

• Cometary late veneer: D/H too high?

• Giant wet asteroid(s)

• Disk snowstorms! (Kuchner, Youdin & Bate)- Snowfall: 1”/day for 104 years

Images: • Earth,• water world (Liss or

Gibson• comets, ast belt

JUPITER

Formation of Outer Planets- Gas-giants: Jupiter and Saturn

i) Mostly gas (thick gaseous envelop)ii) Large rocky cores

- Ice-giants: Uranus and NeptuneNeed to form at a region where ice is available

Outer planets must have formed at a region where gas and icy solid material stay abundant for the duration of their formation

Disk Lifetime & Location of Snow Line

Core-Accretion Model(Gas-giant Planets)

(Pollack et al. 1996)

• Farther out in the protoplanetary disk where the temperature of the gas is lower, the density of solids is enhanced with rocky and icy planetesimals.

• Such an enhancement of the solid density may cause collisional accumulation of solids and results in runaway growth to a mass of approximately10 Earth-masses in 0.5-1 million years.

• These bodies may accrete gas (equivalent to 100 Earth-masses) from the disk within approximately 6-10 million years and form gas-giant planets.

• The gas collapses and forms a thick envelope.

WATER FROM EMBRYOS

Raymond et al., 2004

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Raymond et al., 2004

Stochasticity in the resulting water budget A large eccentric Jupiter inhibits the delivery of water to the inner S.S.

Chambers, 2001; Raymond et al., 2004

• explains the accretion of a LARGE amount of water

• The accreted water has the D/H ratio similar to that of carbonaceous chondritic origin

• The water accretion occurs DURING the formation of the Earth, NOT in a late veneer phase, in agreement with geochemical modeling

• The accretion of the water is a stochastic event, and therefore explains why not all terrestrial planets had an identical primitive water budget (e.g.Mars)

170 Etxrasolar Planets

-Close-in gaint planets(hot Jupiters)

-Eccentric orbits

-Multi-planet systems

-Planets & binary stars

Planetary System vs

Binary Star System

Until a few years ago, it was generally believed that the collapse of a molecular cloud would result in the formation of a planetary system around a single star, or the formation of a dual-star system with no planets.

Observations imply planets in binaries

Circumbinary Disk

GG Tau (a = 35 AU)Md = 0.2 Solar-mass

Debris Disk

HD 141569 separation ~950 AU

Clampin et al. 2003Krist et al. 2005

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• Approximately 20% of extrasolar planets are in binary or multi-star systems

• Almost all these binaries are wide (250-6500 AU)

• γ Cephei (~ 18.5 AU), GJ 86 (~20 AU), and HD188753 (~12 AU) are binary or multi-star systems with at least one Jupiter-like planet

HD142 (GJ 9002) HD3651 HD13445 (GJ 86)HD19994 HD27442 HD40979HD41004 HD75732 (55 Cnc) HD80606 HD89744HD114762 HD 117176 (70 Vir) HD121504HD137759 HD178911 HD186472 (16 Cyg)HD190360 (GJ 777 A) HD192263 HD195019 HD213240HD217107 HD219449 HD219542HD178911 PSR B1257-20 PSR B1620-26

HD9826 (UpsilonHD22049 (Epsilon

HD120136 (TauHD143761 (Rho

HD222404 (Gamma

Binary and multi-star systems with planets(Haghighipour, 2005)

υ Andromedae

http://mcdonaldobservatory.org/news/releases/2002/1009.html

γ Cephei

1.59 solar-mass0.37-0.75 solar-mass

1.7 Jupiter-mass

Triple-star system HD 188753

Primary1.06 MSun

0.96 MSun

0.67 MSun

Porb = 25.7 years, a = 12.3 AU, e = 0.50

Porb = 156 daysa = 0.67 AU

(Konacki, 2005)

Planet=1.14 Jupiter-massPeriod=3.35 days

Giant Planet FormationCurrent theories of planet formation can explain

formation of planets around single starsCore Accretion

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(Artymowicz and Lubow, 1994)

Stellar Companion Affects the Structure of the Nebula

A stellar companion affects the disk by truncating it to 0.5-0.1 times the semimajor axis of the binary

Stellar Companion Affects the Structure of the Nebula

-Single Solar Mass Star-No stellar companion-20 AU radius

-Equal Mass Binary System-Stars = Solar Mass-Binary Semimajor Axis = 50 AU-Binary Eccentricity = 0.5

Boss (2005)

Thiebault et al (2004)

Stellar Companion Affects the Dynamics of Planetesimals

-Increasing eccentricity-Increasing mutual collisions-Increasing the possibility ofcoalescence/ejection

http://mcdonaldobservatory.org/news/releases/2002/1009.html

γ Cephei

1.59 solar-mass0.37-0.75 solar-mass

1.7 Jupiter-mass

Long-Term Stability

Orbital Parameters of γ Cephei

Semimajor Axis = 18.5 ± 1.1 AU Semimajor Axis = 20.3 ± 0.7 AU Eccentricity = 0.361 ± 0.023 Eccentricity = 0.389 ± 0.170

Hatzes et al (2003) Griffin et al (2002)

Numerical Simulation

Binary semimajor eccentricity: 0.2 to 0.65 in steps of 0.05Planet orbital inclination: 0 to 80 degSecondary mass: 0.3 to 0.92 solar-mass

Orbital Stability

Orbit of the Jupiter-size planet is stable for all values of

- binary eccentricity ≤ 0.45

- planet orbital inclination ≤ 60 deg

Planet=Black, Binary=Red

(Haghighipour, 2005)

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F(r) =T

TSun

⎛

⎝ ⎜

⎞

⎠ ⎟

4R

RSun

⎛

⎝ ⎜

⎞

⎠ ⎟

2r

rEarth

⎛

⎝ ⎜

⎞

⎠ ⎟

−2

FSun rEarth( )

F(r) =1

4πL(R,T) r−2 = σ T 4 R2 r−2

Habitable ZoneA habitable zone is a region where an Earth-like planet receives the same amount of radiation as Earth receives from the Sun, and it develops similar habitable conditions as those on the Earth. For a star with luminosity L(R,T), this implies

where

= Star’s brightness

T = Star’s surface temperatureR = Star’s radius r = Radial distance of habitable region from central star

BinaryPeriod = 20750.6579 ± 1568.6 daysSemimajor Axis = 18.5 ± 1.1 AU Eccentricity = 0.361 ± 0.023

Primary SecondaryMass = 1.59 Solar-masses Mass = 0.35-0.75 Solar-massesRadius = 4.66 Solar-radii Radius = 0.5 Solar-radiiTemp = 4900 K Temp = 3500 K Distance = 45 light yearsAge = 3 billion years Planet

Period = 905.574 ± 3.08 daysSemimajor Axis = 2.13 ± 0.05 AUEccentricity = 0.12 ± 0.05Min Mass = 1.7 Jupiter-masses

γ CepheiA Jupiter-like planet in a binary star system

Secondary1.67 Jupiter Mass

1 AU

2.13 AU 18.5 AU

Primary

Habitable Zone

Surface Temperature of primary T = 4900 KRadius of primary R = 4.66 Solar-radii

Habitable zone of γ Cephei : 3.1 AU < r < 3.7 AU

Habitability

The habitable zone of the primary of γ Cephei is UNSTABLE

(Haghighipour, 2006)

1 AU

0.8 AU0.3 AU

Region of Stability of a Terrestrial Planet

2.13 AU 16,17,18 AU

Habitable Zone

Stellar Companion

a = 20, 30, 40 AU

Jupiter

0.5 AU

4 AU

e = 0.0, 0.2, 0.4

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Numerical Simulations

- Binary separation = 20, 30, 40 AU- Binary eccentricity = 0.0, 0.2, 0.4- 120 Embryos randomly distributed from 0.5 to 4 AU- Mass of embryos = 0.01 to 0.1 Earth-mass- Total mass of the disk = 4 Earth-masses- Jupiter at 5 AU- Stochastic => 3 different run for each case

Companion = 1 Solar-mass, Semimajor Axis = 20 AU, Eccentricity = 0

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(Haghighipour & Raymond 2006)

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Companion = 1 Solar-mass, Semimajor Axis = 30 AU, Eccentricity = 0

(Haghighipour & Raymond 2006)

Companion = 1 Solar-mass, Semimajor Axis = 20 AU, Eccentricity = 0.2

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(Haghighipour & Raymond 2006)

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I) The key factor in the amount of water delivered is Jupiter's eccentricity

II) Dynamics of Jupiter is affected by the eccentric orbit of the stellar companion

III) It would be important to understand where giant planets will form in binary systems and to explore whether there is a systematic relation between the binary parameters and the orbit of the outermost giant planet?

Studies of crater densities at sites of known ages (from Apollo samples) give flux data back to ~3.8 Gy ago, and show that the bombardment was ~100 times higher

Cataclysmic LHB (Tera, Ryder, Kring, Cohen, Koeberl..)

Slowly fading LHB (Neukum, Hartman..)

Evidence for HB ~4.0-3.8 Gy ago

- The ages of the rocks collected on the Moon ~3.9-3.8 Gy

-The ages of many basins (impact features > 200km)~3.9-3.8 Gy (Wilhelms, 1987; Ryder, 1994)

Suggests a sudden and short-lived cratering episode ~ 3.9 Gy ago,

(Tera et al. 1974)

•LHB requires a reservoir of small bodies, which have remained stable for ~700 My

•This is possible if there is a change in the orbital structure of the planetary system

NEED TO DELAY THE PROCESS

Planetesimals at farther distances

Planetary eccentricities almost zero(Planet Formation)

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Consider planetesimals only where their dynamical lifetime is not shorter than the gas disk lifetime

Planet positions

Lifetime of planetesimals

Origin of the LHB 1:2 resonance crossing as a function of disk inner edge

1:2 resonance crossing