Module 6: Modelling the Formation of the Solar System. Activity 2: Planetary Evolution. Summary. In this Activity, we will investigate:.
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Module 6: Modelling the Formation of the Solar System
Activity 2:Planetary Evolution
In the previous Activity we learned how protoplanets formed in the Solar Nebula. In this Activity we will have a look in some detail at how they evolved to form the terrestrial (“rocky”) planets in the inner Solar System, and the Jovian (“gas giant”) planets of the outer Solar System.
We will also explore how the planets have evolved since their formation some 4.5 billion years ago.
As we have just seen, close to the Sun this material would have consisted mostly of metal oxides, iron & nickel compounds and silicates - the materials which form the basis of the present day rocky or terrestrial planets - Mercury, Venus, Earth and Mars - and the natural satellites (or moons) of the inner Solar System.
Deformation by big impacts
The rocky planetesimals gradually accreted more material, again due to gravitational attraction:
As the planetesimal grows to planetary size, its interior heats up.
The heating is due to
The continued impact of planetesimals kept the terrestrial protoplanets in a near molten state. As they continued to grow in size, the rocks in the interior of the planets were compressed due to the increase in gravity. The radioactive decay of elements within the rocks also added to their internal heat. If the rate of heating due to these three processes was faster than the rate of cooling, then the planet would heat up.
In the first billion years of the terrestrial planet’s life, its interior is hot enough to melt iron. The dense molten iron sinks to the centre of the planet, and the lighter materials begin to rise towards the surface. This process is called planetary differentiation..
Since gravity is directed towards the centre of a planet, the molten material tried to fall inwards ...
… and so the planets took roughly spherical shapes, then cooled gradually to form brittle outer skins (crusts):
(this is not toscale - the crust would be much thinner than shown here)
Denser (iron-rich) material settles in the centre (core)
Lighter (silicon-rich) material rises towards the surface (mantle)
The idea that planet-sized rocky objects can “melt” due to their own internal energy is pretty surprising, until we remember that the Earth has a molten core, and we see the heat released from the Earth’s still hot mantle in volcanic activity.
There is another particularly clear piece of evidence for this:
if we take a census of Solar System objects, we find that ...
– rocky bodies with diameters greater that 200 km are roughly spherical:
– whereas bodies with diameters less than 130 km are usually irregular:
– which agrees quite well with calculations of how large an accreting object can become before it differentiates.
The process of planetary differentiation also leads to outgassing, whereby internal gasses escape to the surface of the planet and form its outer atmosphere and – in the case of the Earth – oceans. The atmospheres and oceans (of the Earth) are also added to by impacting icy comets in the early history of the Solar System.
Once the planet has differentiated, the interior then gradually cools (though radioactive decay still acts as an internal heating source in the terrestrial planets) and the upper crust solidifies.
In some planets, the mantle and even the core slowly solidifies. The smaller the planet, the quicker it can radiate its internal heat, cool down and finally solidify.
The early Solar System would have contained manyplanetesimals left over from the Solar Nebula.
The planets and natural satellites that we see today in the inner Solar System only represent a fraction of the number of planetesimals and general debris which would initially have been present.
With all this debris around, collisions must have beenquite common:
- some would have caused more accretion, resulting in the growth of planetesimals and protoplanets,
- other more energetic collisions would have broken young planetesimals apart!
As we have seen, the planetesimals which managed to grow large enough to differentiate will have then gradually cooled and formed solid, brittle crusts.
Once solid crusts formed, more impacts with debris in the early Solar System caused extensive cratering:
Cratering evidence exists on all the terrestrial planets, and on all the natural satellites with ancient surfaces.
However we do not see signs of cratering on natural satellites with active (volcanic) or icy surfaces,and we only see limited signs of cratering on Earth
- due to volcanic activity, weathering, extensive plant life, the oceans covering much of the surface, and human activities such as agriculture.
Some spectacular examples of craters do however remain...
The cratering caused cracks in the planet’s crustwhich could be filled up by lava (molten mantle material), heated by radioactive decay, as it welled up through the cracks.
If there was significant liquid water on the young planetit was likely to be present firstly as water vapour.
As the atmosphere cooled, the water would have condensedas rain, filling craters and forming the first oceans.
Long after the crust on a planet’s surface has formed, the mantle may still be hot enough to undergo plastic flow
- that is, move in convective currents like those in water heated in a saucepan on a stove.
If the planet’s interior does not cool down too quickly,the convection currents in the mantle could drag along regions of crust by a few cm per year.
This is what we call plate tectonics, or continental drift, on Earth.
Where plate tectonics occur, as we will see when weinvestigate the Earth, plates can collide with each other,crumpling the crust to form mountain chains and pushing up molten lava to erupt as volcanoes.
Once a planet’s mantle cooled enough to bring its volcanic activity largely to an end, if the planet had an atmosphere, it would then have largely settled down to a long period of gradual weathering, from one or more of:
Which of these happen, and the rate & degree, depended on the atmospheric conditions & circulation patterns on theparticular planet involved.
The ices - such as water, methane and ammonia ices - are made of elements which were much more abundant than the elements which formed the rocky planetesimals.
Once the mass of the rocky and icy protoplanetary cores reached about ten Earth masses, they had enough gravity to begin attracting the surrounding nebula gas, and a gaseous envelope began to form around the cores.
The average speed of gas atoms and molecules depends on the temperature of the gas. In the outer Solar System, even while it was still forming, the temperatures were so low that gas atoms moved very slowly and were easily captured.
Both the core and envelope of the giant planets continued to grow, and as they became more massive, so too did their attractive gravitational force. At some stage, the envelope mass began to increase more rapidly than the core mass and a runaway accretion process followed.
The atmospheres of the giant planets accumulated very rapidly, attracting more and more gas – mostly hydrogen – from the surrounding Solar Nebula.
This means that both the cores and atmospheres of the giant planets must have formed within that time.
It seems likely that more massive Jupiter and Saturn attracted much of the gas in the outer Solar Nebula, leaving less material available to go into producing the atmospheres of the smaller giants Uranus and Neptune.
We’ll return to the subject of their evolution, but first we will spend several Modules studying the planets in more detail, both for their own intrinsic interest and to see what evidence they provide for the model we have been outlining in this Module.
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