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The Solar Atmosphere
Photosphere: visible “surface” of the Sun,
about 500 km thick. Moving outwards
temperature falls from 8000 K to 4500 K and
density of gas rapidly decreases. Granulation
of surface indicates large scale movement of
gas - convection currents transfer heat.
Chromosphere: outside photosphere, about
1500 km thick, temperature rises from
4500 K to 6000 K. Visible as red flash
during solar eclipse.
Corona: starts about 2000 km from the solar
surface, rapid temperature rise to 500,000 K
then slower rise to well over one million K.
Probably heated by electric currents due to
changing magnetic fields.
Total energy in corona is small despite the
high temperatures since the gas is very
tenuous. Becomes the solar wind.
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Meeting the Energy Budget
1370 Watts of solar energy per second hit
every square metre at the distance of the
Earth’s orbit.
Figs. Z13.4 & K11-4
This covers 4p x (Sun to Earth distance)2
square metres of area = 2.81 x 1023 m2
(or ~ 550 million times Earth’s surface area)
Hence Sun’s total radiative energy output
= 3.85 x 1026 Watts
(or ~ ten thousand million million power stations)
Source of energy cannot be chemical since
Sun is too small - it would only burn for a
few thousand years.
Can convert a little mass to a lot of energy:
e = m  c2
(where c2 = 90,000,000,000,000,000)
Source is nuclear energy.
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Luminosity (energy output) of a Star
Need energy reaching Earth and distance
from Earth, just as for Sun.
Parallax (triangulation using Earth’s orbit
as a baseline) gives distance.
Surface temperature of a star - measure
spectrum and fit a black body curve.
Plot luminosity against temperature the Hertzsprung-Russell (H-R)
Diagram
Figs. Z14.17 & K12-17
Weighing Stars - binary systems
Visual binaries :
Figs. Z14.22 & K12-23
need to see the orbits and measure the
distance to Earth; gives the mass.
Spectroscopic binaries :
Figs. Z14.24 & K12-27
Doppler shift gives velocities and orbital
period which can be used to find the mass.
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Eclipsing binaries: measure orbital motion
from variation of apparent brightness.
Plot luminosity against mass (scale Sun = 1)
Figs. Z14.27 & K12-34
Now have a mass-luminosity diagram for
Main Sequence stars.
This tells us that the heavier they are, the
greater is their energy output rate i.e.
Luminosity  (mass  mass  mass  mass)
but
Total energy available  mass
Therefore, low mass stars live longer.
So where does all this energy come from?
n.b.  means “is proportional to”
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Star Birth
Giant molecular gas cloud condenses.
Gas looses gravitational energy and gains
heat energy.
Cloud increases in density and mass, attracts
more gas and warms up more - a protostar.
Eventually the temperature at the centre of
reaches 8,000,000 K and hydrogen “burning”
begins - a star is born.
Figs. Z16.9 & K13-13
Proton-proton chain: fusion of hydrogen to
form helium.
Figs. Z13.15 & K11-7
Each p-p chain provides so little energy that
140 thousand million million kilogrammes of
matter must be converted per year to keep
the Sun shining (equivalent to the mass of our
Moon every 500 thousand years).
Hydrogen is converted to helium in the core.
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The Proton-Proton Chain
neutrino
p  p  p n  e 




positron
hydrogen
nucleus
p n  p  p p n 




deuterium
nucleus

gamma
ray



p p n p p n 
p p nn  p  p  




helium
nucleus
For the Sun, hydrogen in core lasts about
10 billion (10 thousand million) years; we
have just over 5 billion years to go.
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The Inverse Square Law: if we double our
distance from a star then we receive one
quarter of the amount of heat and light - try
stretching your hand out to a light bulb and
gradually backing away; the temperature
drops quickly.
The Hertsprung-Russell Diagram: our Sun is
a very average star. Stars above it are brighter
and stars to the left of it are hotter.