Download Birth of the Solar System

FORMING THE PLANETS
GLG-190 - The Planets
Chapter 16
LECTURE OUTLINE
 Objectives
 General
features of
Solar System
 Historical Models
 Modern Synthesis
OBJECTIVES
Summarize general features of
Solar System
 Any model for formation of
the solar system must explain
large number of observations


Orbital motions, angular
momentum distribution, ages,
sizes and densities, distribution
of small bodies, moons,
features in meteorites,
differentiation, atmospheres,
surface features
ORBITAL MOTIONS
Orbits of planets are roughly circular and coplanar (close to
Sun’s equatorial plane)
 Planets orbit in same direction as Sun rotates (prograde)
 Most planets rotate in same direction as they orbit with axial
tilts less than 30 (Venus and Uranus are exceptions)
 Separation between orbits increases with distance from Sun
 Multitude of smaller bodies (asteroids, KBOs, comets) with
more eccentric and inclined orbits

ANGULAR MOMENTUM DISTRIBUTION

Odd distribution in Solar System…



Sun with >99% of Solar System’s mass
Planets comprising 0.14% of mass have
99.7% of angular momentum
Collapsing gas cloud will spin faster due
to conservation of angular momentum



Angular momentum will concentrated where mass is
concentrated (star at center of spinning cloud)
If mass distribution changes, rate of rotation will
change (skater changes distribution by moving in
arms)
After collapse of protostellar cloud, resulting star
would be spinning very fast
PLANET SIZES AND DENSITIES

Inner Solar System (terrestrial
planets)





Small, dense, rocky
Composed of metal and rock
Volatile (water, gas) poor
Relatively thin atmospheres
Outer Solar System (giant
planets)




Large, less dense, gaseous
Composed of gases and ices
Volatile rich
Massive atmospheres
DISTRIBUTION OF SMALL BODIES


Countless minor planets in Asteroid Belt between Mars and Jupiter (2.1
to 3.3 AU)
Multitude of small bodies (KBOs) in flattened disk of Kuiper Belt beyond
Neptune (35 to 50 AU)



Short-period comets
Many KBOS have moons
Trillions of objects (long-period comets) in spherical Oort Cloud around
Sun located at 10,000 AU
MOONS & PLANETARY RINGS

Moons




Composed of rock and ice
Larger moons in prograde orbits
in plane aligned with planet’s
equator; most tidally locked
Smaller irregular satellites in
retrograde orbits with high
inclinations and/or
eccentricities
Rings around giant planets (in
equatorial planes)


Ring particles have prograde
orbits
Rings located inside orbits of
sizeable moons (inside Roche
limit)
METEORITES


Much diversity in meteorite compositions
Evidence of heating and cooling events
CAIs indicate very high temperatures
 Chondrules were flash melted and rapidly cooled
 Achondrites indicate melting in parent bodies (early heat source)



Preservation of presolar grains  parts of protosolar nebula were
relatively cool
Evidence of mixing of high and low temperature materials (supported
by oxygen isotope results)
METEORITE AND ROCK AGES

Meteorites
Pb-Pb isochron dating of CAIs in
meteorites  4.56 Ga age for oldest
Solar System materials (top)
 Chondrules solidified only few million
years after CAIs
 Narrow age range for meteorites (cluster
around 4.55 Ga)  rapid formation


Rocks
Lunar rocks typically 3 to 4.4 Ga
 Terrestrial rocks all < 4 Ga
 Martian meteorites mostly 150 to 1500
Ma; one sample > 4 Ga


Pb-Pb isochron age for Earth is 4.55
Ga (bottom)
PLANETARY DIFFERENTIATION



All planets, many moons, and meteorites show evidence of
differentiation
Planets have denser cores surrounded by less dense layers
Iron meteorites  cores of differentiated parent bodies
ATMOSPHERES

Terrestrial planets







Thin secondary atmospheres
Depleted in H and He relative to
Sun
Dominated by CO2 and N2
O2 in Earth’s atmosphere
Thick atmosphere of Venus
Faint young Sun Paradox
Giant planets



Very thick primary atmospheres
H and He abundances roughly
similar to solar
Some modification by internal
processes (e.g., He rain)
Terrestrial planet atmospheres
CRATERING AND VOLCANISM

Cratering

Variable surface ages based on crater
densities




Some surfaces saturated  ancient (e.g., Callisto)
Many bodies show multiple ages (e.g., Mars,
Ganymede, Moon)
Surface of Venus essentially uniform age
Some bodies essentially lack craters  very young
(Earth, much of Enceladus, Io)
Present cratering rates too low to produce
high densities  higher impact rates in past
 Evidence of Late Heavy Bombardment at
about 3.8 Ga


Volcanism
Silicate volcanism on all inner Solar System
bodies and Io
 Cryovolcanism on some Moons of outer Solar
System (e.g., Europa, Enceladus, Triton)

A LITTLE HISTORY
Emanuel
Swedenborg
(1688-1772)

Pierre-Simon
Laplace
(1724-1804)
Viktor Safranov
(19175-1999)
George Wetherill
(1925-2006)
Many models proposed to explain features


Immanuel Kant
(1724-1804)
Near collision of stars, capture of planets by Sun, etc.
Processors to modern nebular hypothesis proposed by Swedenborg
(1734), Kant (1755), Laplace (1796)
Rotating sphere of gas flattens into spinning disk
Contraction of disk causes disk to spin faster
Gas in disk forms planets
 Angular momentum problem: Sun should be spinning rapidly




Solar Nebular Disk Model (SNDM): Safronov (1969) and Wetherill (1970s)

Planetesimals rather than gas accretion
WHAT WE WANT IN A FORMATION MODEL
 Incorporate
key observations into consistent model of
solar system formation
 Use astronomical observations to establish events leading
to protoplanetary disk
A GENERAL MODEL

Collapse of cold interstellar
molecular clouds


Rotation rate increases, cloud
flattens into spinning disk



Collapse triggered by shock
Enveloped in cocoon of gas and
dust
Disk surrounded by other stars
Accretion of material within
disk  protoplanets
Temperature controls materials
condensation
 Remaining gas swept away by TTauri wind



Protoplanets merge to make
planets
Giant collisions in final stages
GIANT MOLECULAR CLOUDS
Huge clouds of gas and dust (mass
103-107 times Msun)
 Mixture of materials with different
origins
 Densest parts collapse into stars
 Largest stars emit huge amounts of
UV





Pushes gases away leaving cavity
Causes gas to glow (HII region)
Large stars have short lifetimes
Smaller stars also form
Age more slowly
Evolution influenced by presence of
larger stars
 Incorporate materials from older and
larger stars
 Develop protoplanetary disks


STAR-DISK FORMATION IN HII REGIONS



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
Star-formation triggered by
compression of cold gas around HII
region (right)
Denser clumps appear (evaporating
gaseous globules, “EGGs”) as ionization
moves into cold surrounding gas
Exposed EGGs are shaped by radiation
from nearby large stars  tear-drop
Proplyds (protoplanetary disks) form
within EEGs (new star at center)
Gas is completely stripped away yielding
“naked” disk in hot tenuous interior of
HII region
Older material may be added to
protoplanetary disk from nearby
supernova explosions and large stars
Massive
stars
Molecular
cloud
HII region
New stars form
Erosion of cloud
Model of Hester & Desch, 2005
EAGLE NEBULA
“Pillars of
Creation”
EGGS IN THE EAGLE NEBULA
EGG TRANSFORMING TO PROPLYD
(jet from hidden YSO)
Trifid Nebula
PROPLYDS IN THE ORION NEBULA
Forming a disk
BARE PROTOPLANETARY DISKS IN ORION
BIPOLAR JETS

Material falling into star at
center produces bipolar flows
 material ejected into
interstellar medium
Hubble image of the
Egg Nebula
SUPERNOVAE
Disk is near many very large
stars, which are common in
HII regions
 Large stars have very short
life spans (typically 3-30
Myr) and explode in
supernovae
 Supernovae form…




Elements heavier than iron
(up to uranium)
Short-lived radionuclides from
supernovae (60Fe  60Ni, t½ =
1.5 Ma)
Material ejected into space
and into protosolar nebula
EXTRASOLAR GRAINS

Grains form around giant stars
and ejected into interstellar
space
High temperature materials:
diamond, silicon carbide (upper
right), graphite, corundum,
spinel, etc.
 Have non-solar isotopic
compositions (diagram at right)


Grains captured by protosolar
nebula
SiC
EVOLUTION OF PROTOPLANETARY DISKS
SOLVING THE ANGULAR MOMENTUM
PROBLEM


Recall, conservation of angular
momentum requires…
 Rotating disk forms with most
momentum in proto-Sun
 BUT most angular momentum
held by outer planets
Sun lost most of its angular
momentum  transferred to outer
planets by magnetic braking
 Magnetic lines of force sweep
through nebula
 Charged particles “dragged” along
field lines  transfer of
momentum from Sun to nebula
MAKING GRAINS: CONDENSATION
DISK TEMPERATURE VARIATIONS

Dominated by star in center
Near to the central proto-sun, the nebular
temperature will be very high  no solids
can condense
 Farther away from proto-sun, temperatures
fall off


Condensation Sequence




Beyond 0.2 AU, temperature below 2,000 K
 metals and oxides (corundum, spinel)
condense
At 0.5 AU, temperature below 1,5000 K 
silicate minerals (olivine, pyroxene)
condense
Beyond 5 AU, temperature below 200 K
 ices condense
Temperature (distance) controlled
sequence of chemical condensation
 correctly predicts basic chemical
make-up of planets
INTERPLANETARY DUST PARTICLE
Example of a particle aggregate…
ACCRETION STAGES
Formation of planetesimals (10m to 1000 km) (10,000
year time scale)
 Growth of planetesimals by collisions/intersecting
orbits (million year time scale)
 Formation of planetary “embryos” with masses like
Moon and Mars (million year time scale)

Embryos collide to form planets
 Earth-Moon system result of such collision
 Many other features explained by collisions


About 100 Myr between initial condensation and
formation of Earth-Moon system
ACCRETION PROCESSES




Initially, dust sized
particles aggregate
by sticking and
coagulation (gravity
does not play
significant role)
Over 20,000 years,
particles grow into
planetesimals
(diameters 1-10s
km)  gravitational
attraction becomes
dominant process
With planetary diameters of 1000s to 10,000s km, gravitational attraction
sufficient to sweep up gases in orbital vicinity
Entire process probably takes few million years
FORMATION OF TERRESTRIAL PLANETS
ACCRETION IN INNER SOLAR SYSTEM


Simulation of final stage of
growth of terrestrial planets
Note significant orbital changes
 mixed material from
different distances
Simulation source: http://casa.colorado.edu/~raymonsn/graphics.html


Begins with planetary embryos with
masses between that of Moon and
Mars
Ends with planets similar to inner
planets of Solar System
WHY NO PLANETS IN ASTEROID BELT?

Resonances associated with
giant planets remove nearby
protoplanets
Kirkwood gaps indicate present
positions of resonances
 Location of resonances will shift
when the orbits of giant planets
shift



Gravitational interactions
perturb other protoplanets into
resonances until all are lost
Planets survive within 2 AU of
the Sun where there are no
resonances
FORMATION OF GIANT PLANETS
“FROST” LINE



Also called “snow line”
Inner solar system too hot for
water to freeze
At sufficient distance (3 AU),
water will freeze out  larger
planets form (planets then can
retain H and He)
FORMATION OF JOVIAN PLANETS

Formation of gas giants
favored just outside “frost
line” at 3 AU
Accretion builds large cores of
rock and ice  Jupiter and
Saturn have large core (10-15
MEarth)
 Disks form around cores due to
gravitational attraction…
 Gravitational capture of H, He,
and dust from nebula to form
“mini” disks


Formation retarded in outer
regions (10 AU)  not as many
planetesimals and as much gas
available
Swirling dust
and gas around
AB Aur (100x
bigger than
our solar
system)
CLEARING THE DISK: T-TAURI WINDS

When density of gas is sufficient for hydrogen ignition
(fusion), protostar becomes luminous object

Strong solar winds clear nebula of gas and dust (may strip away
terrestrial atmospheres)  end of accretion
PROTOPLANETARY DISK BRIGHTNESS
Brightness reflects both size and nature of disk
 Disk brightness decreases with increasing star age 
evidence for accretion
 Very bright disks uncommon after 50 my

PLANETARY DIFFERENTIATION



Separation of material based upon
density
Requires high temperatures to
allow internal movement (flow) of
material
Large bodies heat more effectively
 low surface area to volume ratio


Heat from decay of long-lived
radioisotopes (235U, 238U, 232Th, 40K)
Evidence that even smaller
meteorite parent bodies
differentiated

Heat producing by radioactive decay of
26Al (decays into stable 26Mg with halflife of 0.73 Ma)
CORE FORMATION




Several models

Instantaneous

Continuous (accretion of Earth)

Core merging (Moon-forming impact)
Time of core formation can be
determined using Hf-W isotopes

Earth 30 Ma after formation

Mars 15 Ma after formation
Cores probably start out completely
molten
Iron-loving “siderophile” elements
(Au, PGE) sequestered into core
(stay in metal alloy)
GIANT IMPACTS
Last stages of accretion
 Collisions of large
planetesimals and planets
 Impacts proposed to
produce…


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
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

Mercury’s large core
Retrograde rotation of Venus?
Earth’s Moon
Crustal dichotomy on Mars
Tilt of axis of Uranus
Pluto-Charon system
Impacts also help remove
primary atmospheres
PLANETARY MIGRATION



Interaction of large
planets with abundant
planetesimals in outer
solar system
Inner giant planets can
migrate inward (Jupiter)
 “hot Jupiters” seen in
extrasolar systems
(right)
Movements only stops
when the supply of
planetesimals
disappears
Simulation source: http://casa.colorado.edu/~raymonsn/graphics.html
PRIMARY ATMOSPHERES



Composed mostly of light
gases accreted during initial
formation

Same mixture as found in Sun
and Jupiter (gases directly from
solar nebula)

Roughly 94.2% H, 5.7% He
(everything else <0.1%)
Lost from terrestrial planets

During accretion

Later by thermal escape and
solar wind ablation
Retained around giant planets

High gravity, low temperatures
Bodies retain all gases with lines running
below their plotted positions
SECONDARY ATMOSPHERES

Venus, Earth and Mars have secondary atmospheres


Produced impact degassing of volatile-rich asteroids and
volcanism (very minor role for comets)
Presumed to have similar initial compositions


Rough proportions of main components: 83% H2O, 17% CO2 and
0.5% N2
Atmospheres of Venus and Earth probably thicker than Mars
GOLDILOCKS PARADOX
Why have Venus, Earth, and Mars
ended up with different
atmospheres? H2O is key!
 Earth (“just right”)





Temperature allows liquid water
oceans (most H2O removed from
atmosphere)
CO2 dissolves in water (carbonic acid)
and reacts with rocks  limestone
Removing H2O and CO2 from
atmosphere  N2 dominates (78%)
Further modified by life (3.5 Ga),
which consumes CO2 to produce
oxygen (21%)

Venus (“too hot”)





Too hot for liquid water  no place for CO2 to
dissolve
Hot “wet” greenhouse (CO2 and H2O)
Water in upper atmosphere disassociates to form
O (reacts with surface rocks) and H (lost to space)
CO2 retained as main component (97%) in thick
atmosphere  runaway “dry” greenhouse effect
Mars (“too cold”)



Liquid water present soon after formation, but
gradually freezes out  decreases greenhouse
effect
Most CO2 freezes out, remainder dominates
(95%) thin atmosphere (0.007 bar)
Magnetic field shuts down early  atmospheric
loss by solar wind stripping?