Download Formation of the Solar System

Formation of the
Solar System
The Age of the Solar System
We can estimate the age of the Solar System by looking at
radioactive isotopes. These are unstable forms of elements that
produce energy by splitting apart (i.e., fission).
The radioactivity of an isotope is characterized by its half-life –
the time it takes for half of the parent to decay into its daughter
element. By measuring the ratio of the parent to daughter, one
can estimate how long the material has been around.
Radioactive Elements
Isotope
#
#
protons neutrons
Daughter
Half-life (years)
Rubidium-87
37
50
Strontium-87
47,000,000,000
Uranium-238
92
146
Lead-206
4,510,000,000
Uranium-235
92
143
Lead-207
710,000,000
Potassium-40
19
21
Argon-40
1,280,000,000
Aluminum-26
13
13
Magnesium-26
730,000
Carbon-14
6
8
Nitrogen-14
5,730
Each of these isotopes spontaneously decays into its daughter.
In each case, the daughter weighs less than the parent – energy
is produced.
Age of the Solar System
When rocks are molten, heavier elements (such as uranium) will
separate out from other elements. (In liquids, dense things sink,
light things rise.) Once the rocks solidify, the material can no
longer differentiate. Lighter elements (made from radioactive
decay) stay in the same location as they form.
• On Earth, most old rocks have ages of 3 billion years
• The oldest asteroids have ages of 4.5 billion years
• Rocks from the “plains” on the Moon have ages of about 3
billion years. The oldest Moon rocks have ages of 4.5
billion years.
The solar system is therefore 4.5 billion years old.
Keys to Solar System Formation
Any theory for the formation of the Solar System must explain
 The flatness of the Solar System, and orbital similarities
 The separation of Terrestrial and Jovian planets
 The decrease in planet densities with distance from the Sun
 Bode’s Law
Star/Planet Formation
The story of planet formation is in large part, the story of star
formation. Inside dense interstellar clouds of gas and dust, the
temperature is just a few degrees above absolute zero. Since
the temperature is so low, there is no gas pressure to resist
gravity. The cloud collapses.
Initial Collapse
Dark clouds are much
denser in their center than
on the outside, so their inner
regions collapse first.
Also, since the clouds are
lumpy to begin with, the
collapse process causes the
clouds to fragment.
Each fragment is a protostar.
Formation of the Solar Nebula
In a large, slowly rotating cloud of cold gas
 Self gravity begins to
collapse the cloud
 As the cloud gets smaller, it
begins to rotate faster, due to
conservation of angular
momentum.
 Centripetal force prevents
gas from collapsing in the
plane of rotation
 Gas falling from the top
collides with gas falling
from the bottom and sticks
together in the ecliptic plane
Formation of the Solar Nebula
In the flat solar nebula
 The densest region (the
center) becomes the Sun.
Friction in the disk causes
the Sun to accrete matter and
grow in mass. Eventually,
fusion occurs.
 Atoms orbiting in the disk
bump together and form
molecules, such as water.
Droplets of these molecules
stick together to form
planetesimals.
Formation of the Solar Nebula
Planetesimals grow …
 Over time, the planetesimals
grow as more molecules
condense out of the nebula
 Differential rotation (due to
Kepler’s laws) cause
particles in similar orbits to
meet up. They stick together
forming a bigger body.
 The bigger the body, the
greater its gravity, and the
more attraction it has for
other bodies. Protoplanets
form.
Formation of the Solar Nebula
Material begins to evaporate
 While protoplanets are
forming, the Sun’s
luminosity is growing, first
due to gravitational
contraction, then due to
nuclear ignition.
 Regions of the nebula close
to the Sun will get hot; the
outer regions will stay cool.
In the hot regions, light
elements will evaporate;
only heavy elements will
condense out of the nebula
Temperature of the Solar Nebula
Inside the orbit of the Earth, only metals can condense out of
the solar nebula. Rocky (silicates) can condense near Mars.
In the outer solar system, water and ammonia ice can survive.
Radiation Pressure and the Solar Wind
Two other processes are also important for driving light gases
from the inner part of the solar system.
Radiation pressure: Photons act like
particles and push whatever particles
and dust they run into.
Solar wind: The Sun constantly ejects
(a little) hydrogen and helium into
space. This solar wind pushes
whatever gas and dust it runs into.
The Pre-Main Sequence Sun
As the Sun formed, it generated
its energy via gravitational
contraction. During this time, it
was a lot brighter than it is today.
The radiation pressure in the inner
solar system was greater.
In addition, due to conservation
of angular momentum, the young
Sun was also spinning faster than
it is today. This caused the solar
wind to be stronger.
The Pre-Main Sequence Sun
As the Sun formed, it generated
its energy via gravitational
contraction. During this time, it
was a lot brighter than it is today.
The radiation pressure in the inner
solar system was greater.
In addition, due to conservation
of angular momentum, the young
Sun was also spinning faster than
it is today. This caused the solar
wind to be stronger.
Radiation pressure and the solar wind
blew out the light material from the
inner part of the solar system.
The Protoplanetary Disk
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Accretion
Once the major bodies of the solar system were formed, most of
the remaining debris was either ejected out of the solar system
or accreted onto other bodies by gravitational encounters.
Accretion
Once the major bodies of the solar system were formed, most of
the remaining debris was either ejected out of the solar system
or accreted onto other bodies by gravitational encounters.
QuickTime™ and a
YUV420 codec decompressor
are needed to see this picture.
Accretion
Once the major bodies of the solar system were formed, most of
the remaining debris was either ejected out of the solar system
or accreted onto other bodies by gravitational encounters.
QuickTime™ and a
Sorenson Video 3 decompressor
are needed to see this picture.
Accretion
Once the major bodies of the solar system were formed, most of
the remaining debris was either ejected out of the solar system
or accreted onto other bodies by gravitational encounters.
QuickTime™ and a
TIFF (Uncompressed) decompressor
are needed to see this picture.
Unless a body is wellseparated from everything
else, or its orbit is in a
resonance, its orbit will be
chaotic. Eventually, it
will either crash into
something, or leave the
solar system completely.
Accretion
Once the major bodies of the solar system were formed, most of
the remaining debris was either ejected out of the solar system
or accreted onto other bodies by gravitational encounters.
Formation of the Solar System
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From interstellar cloud to planetary system
Observations of Protostellar Disks
Our technology is just beginning to be able to resolve
the proto-planetary disks around stars.
Observations of Protostellar Disks
Our technology is just beginning to be able to resolve
the proto-planetary disks around stars.
Evolution of Terrestrial Planets
After the condensation and accretion phases of planet formation,
terrestrial bodies can go through 4 different stages of evolution.
(The rates of evolution can vary greatly.)
 Differentiation – in a molten planet, heavy materials sink
Differentiation
Early in the history of the solar system, planets would be
molten due to
Continuous accretion of left over
material from the solar system
formation.
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Energy from the fission of
radioactive isotopes.
Evolution of Terrestrial Planets
After the condensation and accretion phases of planet formation,
terrestrial bodies can go through 4 different stages of evolution.
(The rates of evolution can vary greatly.)
 Differentiation – in a molten planet, heavy materials sink
 Cratering – left over bodies impact the planet’s surface
Evolution of Terrestrial Planets
After the condensation and accretion phases of planet formation,
terrestrial bodies can go through 4 different stages of evolution.
(The rates of evolution can vary greatly.)
 Differentiation – in a molten planet, heavy materials sink
 Cratering – left over bodies impact the planet’s surface
 Flooding – water, lava, and gases trapped inside the planet
come to the surface and cover the terrain.
Evolution of Terrestrial Planets
After the condensation and accretion phases of planet formation,
terrestrial bodies can go through 4 different stages of evolution.
(The rates of evolution can vary greatly.)
 Differentiation – in a molten planet, heavy materials sink
 Cratering – left over bodies impact the planet’s surface
 Flooding – water, lava, and gases trapped inside the planet
come to the surface and cover the terrain.
 Erosion – surface features are destroyed due to running
water, atmosphere, plate tectonics, and geologic motions