Download Nebula Theory

Planet Formation
From Dust to Planetesimals
Stopping time
Coagulation of dust grains forms millimeterand centimeter-sized objects
(Dominik and Tielens, 1997)
MAGIC
Physical Characteristics
Drag forces
Strong -> Week
Chemical binding
Strong -> Week
Surface gravity
Week -> Strong
Formation of Terrestrial Planets
Runaway Growth
Gravitational interaction
causes collisions among
planetesimals and results in
the formation of Mercuryto Mars-sized objects
(planetary embryo)
Collisional coalescence of cm-sized particles
results in the formation of larger objects and
eventually planetesimals (km-sized).
Planetary embryos are
formed in ~10,000 y,
separated by a few
mutual Hill radii.
Accretion of embryos is a
local process.
Ida and Makino (1993)
Kokubo and Ida (1995, 1996, 1998)
1
Final Stage
What if the embryos existed also in the asteroid belt?
Giant impacts among high velocity embryos that result
in terrestrial planets in ~100 million years.
What if the embryos existed also in the asteroid belt?
Water & Earth
- Current location of Earth
too close to the Sun to
retain water
-The icy bodies appear
at distances of 4.0 AU
and larger
-Earth must have acquired
its water from larger
distances
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)?
D/H (x 10-6)
Halley
260-350
Hyakutake 280-300
Hale Bopp 250-410
Courtesy of F. Robert
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
2
WATER FROM ASTEROIDS
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 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
Formation of Outer Planets
- Gas-giants: Jupiter and Saturn
i) Mostly gas (thick gaseous envelop)
ii) Large rocky cores
JUPITER
- Ice-giants: Uranus and Neptune
Need 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
WATER FROM EMBRYOS
(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.
Raymond et al., 2004
3
Stochasticity in the resulting water budget
Raymond et al., 2004
A large eccentric Jupiter inhibits the delivery of water to the inner S.S.
Chambers, 2001; Raymond et al., 2004
170 Etxrasolar Planets
• 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)
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.
-Close-in gaint planets
(hot Jupiters)
-Eccentric orbits
-Multi-planet systems
-Planets & binary stars
Observations imply planets in binaries
Circumbinary Disk
Debris Disk
GG Tau (a = 35 AU)
Md = 0.2 Solar-mass
HD 141569
separation ~950 AU
Krist et al. 2005
Clampin et al. 2003
4
• 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
Binary and multi-star systems with planets
6H
ˆ
HD142 (GJ 9002)
HD19994
HD41004
HD114762
HD137759
HD190360 (GJ 777 A)
HD217107
HD178911
(Haghighipour, 2005)
6H
ˆ
HD3651
HD22049 (Epsilon(ˆ
P
HD75732 (55 Cnc)
HD 117176 (70 Vir)
HD143761 (Rho&ˆ
I
HD192263
HD219449
PSR B1257-20
6H
ˆ
HD9826 (Upsilon$„K
HD27442
HD80606
HD120136 (Tau%††
HD178911
HD195019
HD219542
PSR B1620-26
6H
ˆ
HD13445 (GJ 86)
HD40979
HD89744
HD121504
HD186472 (16 Cyg)
HD213240
HD222404 (Gamma&L‡O
L
P
γ Cephei
0.37-0.75 solar-mass
1.59 solar-mass
υ Andromedae
1.7 Jupiter-mass
http://mcdonaldobservatory.org/news/releases/2002/1009.html
Triple-star system HD 188753
Giant Planet Formation
(Konacki, 2005)
Current theories of planet formation can explain
formation of planets around single stars
0.96 MSun
Primary
1.06 MSun
Core Accretion
0.67 MSun
Porb = 156 days
a = 0.67 AU
Planet=1.14 Jupiter-mass
Period=3.35 days
Porb = 25.7 years, a = 12.3 AU, e = 0.50
5
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)
(Artymowicz and Lubow, 1994)
Stellar Companion Affects the
Dynamics of Planetesimals
γ Cephei
-Increasing eccentricity
-Increasing mutual collisions
-Increasing the possibility of
coalescence/ejection
0.37-0.75 solar-mass
1.59 solar-mass
Thiebault et al (2004)
1.7 Jupiter-mass
http://mcdonaldobservatory.org/news/releases/2002/1009.html
Long-Term Stability
Planet=Black, Binary=Red
Orbital Stability
Orbital Parameters of γ Cephei
Semimajor Axis = 18.5 ± 1.1 AU
Eccentricity = 0.361 ± 0.023
Hatzes et al (2003)
Semimajor Axis = 20.3 ± 0.7 AU
Eccentricity = 0.389 ± 0.170
Orbit of the Jupiter-size planet is
stable for all values of
Griffin et al (2002)
- binary eccentricity ≤ 0.45
- planet orbital inclination ≤ 60 deg
Numerical Simulation
Binary semimajor eccentricity: 0.2 to 0.65 in steps of 0.05
Planet orbital inclination: 0 to 80 deg
Secondary mass: 0.3 to 0.92 solar-mass
(Haghighipour, 2005)
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γ Cephei
Habitable Zone
A 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
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where
F(r) =
−2
⎛ T ⎞ ⎛ R ⎞ ⎛ r ⎞
F(r) = ⎜
⎟ ⎜
⎟ ⎜
⎟ FSun (rEarth )
⎝ TSun ⎠ ⎝ RSun ⎠ ⎝ rEarth ⎠
2
1
L(R,T) r −2 = σ T 4 R 2 r−2 = Star’s brightness
4π
T = Star’s surface temperature
R = Star’s radius
r = Radial distance of habitable region from central star
A Jupiter-like planet in a binary star system
Binary
Period = 20750.6579 ± 1568.6 days
Semimajor Axis = 18.5 ± 1.1 AU
Eccentricity = 0.361 ± 0.023
Primary
Mass = 1.59 Solar-masses
Radius = 4.66 Solar-radii
Temp = 4900 K
Distance = 45 light years
Age = 3 billion years
Surface Temperature of primary T = 4900 K
Secondary
Mass = 0.35-0.75 Solar-masses
Radius = 0.5 Solar-radii
Temp = 3500 K
Planet
Period = 905.574 ± 3.08 days
Semimajor Axis = 2.13 ± 0.05 AU
Eccentricity = 0.12 ± 0.05
Min Mass = 1.7 Jupiter-masses
Habitability
Radius of primary R = 4.66 Solar-radii
Habitable zone of γ Cephei : 3.1 AU < r < 3.7 AU
The habitable zone of the primary of γ Cephei is UNSTABLE
(Haghighipour, 2006)
Habitable Zone
2.13 AU
18.5 AU
1 AU
1.67 Jupiter Mass
Secondary
Primary
Region of Stability of a Terrestrial Planet
Habitable Zone
2.13 AU
a = 20, 30, 40 AU
e = 0.0, 0.2, 0.4
16,17,18 AU
0.5 AU
1 AU
4 AU
Stellar Companion
0.8 AU
Jupiter
0.3 AU
7
Companion = 1 Solar-mass, Semimajor Axis = 20 AU, Eccentricity = 0
Numerical Simulations
(Haghighipour & Raymond 2006)
- 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
1
Companion = 1 Solar-mass, Semimajor Axis = 30 AU, Eccentricity = 0
(Haghighipour & Raymond 2006)
2
Companion = 1 Solar-mass, Semimajor Axis = 20 AU, Eccentricity = 0.2
(Haghighipour & Raymond 2006)
2
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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
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)
Cataclysmic LHB
(Tera, Ryder, Kring,
Cohen, Koeberl..)
Suggests a sudden and short-lived cratering
episode ~ 3.9 Gy ago,
Slowly fading LHB
(Neukum, Hartman..)
•LHB requires a reservoir of small bodies, which have remained
stable for ~700 My
(Tera et al. 1974)
NEED TO DELAY THE PROCESS
•This is possible if there is a change in the orbital structure of the
planetary system
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
Lifetime of
planetesimals
Planet positions
Origin of the LHB
1:2 resonance crossing as a function of disk inner edge
1:2 resonance crossing
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