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56 Cards in this Set

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  • Back
p-p chain
Protons fuse one by one to make helium. This reaction is most important in stars with less than 1 solar mass.
triple-alpha process
After all of a star's hydrogen is fused to helium, other nuclear reactions become important.
Helium fuses into carbon in most stars lying above the main sequence in the HR diagram.
CNO cycle
Protons fuse to a carbon nucleus to make nitrogen and then oxygen. The oxygen nucleus splits apart, making helium and carbon. This reaction is important stars with more than 1 solar mass.
ordinary gas pressure
• All ordinary gas consists of many atoms, separated by much empty space. The atoms travel at high speed and often bump into each other.
• If the atoms are in a box, they bump into the walls of the box. Gas has so many atoms that the individual bumps blur together into a steady force on the walls of the box. This steady force is what we call pressure (actually pressure is the force per unit area)
• The pressure exerted by a gas increases as its temperature increases. At higher temperatures the atoms move faster and hit the walls with greater force.
• The pressure exerted by a gas increases as its density increases. If there are more atoms, then the walls are hit more often and the force on the walls is higher.
• All main sequence stars are supported by ordinary gas pressure.
electron degeneracy pressure
• If a gas is compressed enough, the electrons surrounding the nuclei of the atoms begin to touch. The electrons form a cloud that fills all available space. This is called electron degeneracy.
• Since the electrons fill all available space, a degenerate gas can only be compressed by compressing the electrons themselves. Electrons resist compression and their resistance causes pressure, called electron degeneracy pressure.
• Electron degeneracy pressure depends only on the density of a gas, not on its temperature.
• The densities of white dwarfs are so high (106 gm/cm3) that their electrons are degenerate. All white dwarfs are supported by electron degeneracy pressure.
neutron degeneracy pressure
• If electrons are compressed enough, they will combine with protons to form neutrons, forming a cloud of touching neutrons. This is called neutron degeneracy,
• Since the neutrons fill all available space, a gas of degenerate neutrons can only be compressed by compressing the neutrons themselves. Neutrons resist compression and their resistance causes pressure, called neutron degeneracy pressure.
• The densities in neutron stars are so high (1015 gm/cm3) that their neutrons are degenerate. All neutron stars are supported by neutron degeneracy pressure.
• Neutrons can be compressed, but if they are compressed too much, they will collapse. If this happens in a neutron star, the neutron star will collapse to a black hole.
p + p -> 2H + e+ + n
Meaning of the symbols:
• 2H is hydrogen with an extra neutron. In this reaction the two protons join together but one changes to a neutron. Hydrogen with an extra neutron is called deuterium.
• e+ is a positron. Positrons are identical to electrons except that they have a positive charge. Positrons are the anti-particles of electrons, which means that if a positron collides with an electron, both are annihilated, producing two gamma rays (two photons).
• n is a neutrino. Neutrinos are ghostly particles with no electric charge, almost no mass, that do not interact strongly with other particles.
Energy is released by this reaction in two ways
• The deuterium and the positron are emitted with high velocities (high temperature!).
• The positron annihilates a nearby electron, producing 2 gamma rays (electromagnetic radiation).
• The neutrino also carries away some energy, but the energy generally cannot be used.
E=mc2
• This equation applies to all reactions that produce energy, even chemical reactions, but most reactions do not produce enough energy for the loss in mass to be measured. Nuclear reactions produce so much energy that the loss in mass is easily measured, so Einstein's equation is often associated with nuclear energy.
• Note that Einstein's equation does not say which nuclear reactions will take place nor how much energy the reactions will produce.
hydrostatic equilibrium
• Gravity pulls material towards the center of the star and pressure pushes material towards the surface.
• Pressure and gravity exactly balance each other everywhere in most stars. If they did not, the star would shrink or expand. (This is exactly what happens in pulsating stars.)
• When pressure and gravity are in balance, it's called "pressure equilibrium" or "hydrostatic equilibrium," and the pressure in the star is particularly easy to calculate.
• An important result of the calculations is:
The pressure in a star increases towards its center.
conduction
the transport of energy between objects in direct contact. Conduction is important only in degenerate gasses. It dominates heat transport in white dwarfs and neutron stars.
convection
the transport of energy by the motion of hot gas. Convection is an efficient way to transport energy. It is important in some parts of all main sequence and red giant stars.
radiation
the transport of energy by light (rember that photons carry energy). Radiative transport is generally inefficient because it is easy to absorb light (stars are opaque), but even so, radiative transport is also important in some parts of most stars on or above the main sequence.
solar mass of the sun
1
core temperature of the sun
15 x 106 K
core density of the sun
150 gm/cm3
p-p chain
Protons fuse one by one to make helium. This reaction is most important in stars with less than 1 solar mass.
triple-alpha process
After all of a star's hydrogen is fused to helium, other nuclear reactions become important.
Helium fuses into carbon in most stars lying above the main sequence in the HR diagram.
CNO cycle
Protons fuse to a carbon nucleus to make nitrogen and then oxygen. The oxygen nucleus splits apart, making helium and carbon. This reaction is important stars with more than 1 solar mass.
ordinary gas pressure
• All ordinary gas consists of many atoms, separated by much empty space. The atoms travel at high speed and often bump into each other.
• If the atoms are in a box, they bump into the walls of the box. Gas has so many atoms that the individual bumps blur together into a steady force on the walls of the box. This steady force is what we call pressure (actually pressure is the force per unit area)
• The pressure exerted by a gas increases as its temperature increases. At higher temperatures the atoms move faster and hit the walls with greater force.
• The pressure exerted by a gas increases as its density increases. If there are more atoms, then the walls are hit more often and the force on the walls is higher.
• All main sequence stars are supported by ordinary gas pressure.
electron degeneracy pressure
• If a gas is compressed enough, the electrons surrounding the nuclei of the atoms begin to touch. The electrons form a cloud that fills all available space. This is called electron degeneracy.
• Since the electrons fill all available space, a degenerate gas can only be compressed by compressing the electrons themselves. Electrons resist compression and their resistance causes pressure, called electron degeneracy pressure.
• Electron degeneracy pressure depends only on the density of a gas, not on its temperature.
• The densities of white dwarfs are so high (106 gm/cm3) that their electrons are degenerate. All white dwarfs are supported by electron degeneracy pressure.
neutron degeneracy pressure
• If electrons are compressed enough, they will combine with protons to form neutrons, forming a cloud of touching neutrons. This is called neutron degeneracy,
• Since the neutrons fill all available space, a gas of degenerate neutrons can only be compressed by compressing the neutrons themselves. Neutrons resist compression and their resistance causes pressure, called neutron degeneracy pressure.
• The densities in neutron stars are so high (1015 gm/cm3) that their neutrons are degenerate. All neutron stars are supported by neutron degeneracy pressure.
• Neutrons can be compressed, but if they are compressed too much, they will collapse. If this happens in a neutron star, the neutron star will collapse to a black hole.
p + p -> 2H + e+ + n
Meaning of the symbols:
• 2H is hydrogen with an extra neutron. In this reaction the two protons join together but one changes to a neutron. Hydrogen with an extra neutron is called deuterium.
• e+ is a positron. Positrons are identical to electrons except that they have a positive charge. Positrons are the anti-particles of electrons, which means that if a positron collides with an electron, both are annihilated, producing two gamma rays (two photons).
• n is a neutrino. Neutrinos are ghostly particles with no electric charge, almost no mass, that do not interact strongly with other particles.
Energy is released by this reaction in two ways
• The deuterium and the positron are emitted with high velocities (high temperature!).
• The positron annihilates a nearby electron, producing 2 gamma rays (electromagnetic radiation).
• The neutrino also carries away some energy, but the energy generally cannot be used.
E=mc2
• This equation applies to all reactions that produce energy, even chemical reactions, but most reactions do not produce enough energy for the loss in mass to be measured. Nuclear reactions produce so much energy that the loss in mass is easily measured, so Einstein's equation is often associated with nuclear energy.
• Note that Einstein's equation does not say which nuclear reactions will take place nor how much energy the reactions will produce.
hydrostatic equilibrium
• Gravity pulls material towards the center of the star and pressure pushes material towards the surface.
• Pressure and gravity exactly balance each other everywhere in most stars. If they did not, the star would shrink or expand. (This is exactly what happens in pulsating stars.)
• When pressure and gravity are in balance, it's called "pressure equilibrium" or "hydrostatic equilibrium," and the pressure in the star is particularly easy to calculate.
• An important result of the calculations is:
The pressure in a star increases towards its center.
conduction
the transport of energy between objects in direct contact. Conduction is important only in degenerate gasses. It dominates heat transport in white dwarfs and neutron stars.
convection
the transport of energy by the motion of hot gas. is an efficient way to transport energy. It is important in some parts of all main sequence and red giant stars.
radiation
the transport of energy by light (rember that photons carry energy). Radiative transport is generally inefficient because it is easy to absorb light (stars are opaque), but even so, radiative transport is also important in some parts of most stars on or above the main sequence.
solar mass of the sun
1
core temperature of the sun
15 x 106 K
core density of the sun
150 gm/cm3
• The radius of a black hole is proportional to its mass
• R = 3M, where R is in kilometers and M in solar masses.
age of the sun
4.7 x 10^9
radius of the sun
6.96 x 10^10 cm.
solar granulation
A pattern of bright and dark regions on the surface of the sun

caused by these effects and confirms that the outermost layers of the sun are convective
solar neutrino experiment
Nuclear reactions produce neutrinos in the sun's core, and the neutrinos easily pass through the sun's overlying layers and escape into space. Neutrino telescopes detect only about 1/3 as many neutrinos as the models predict. Originally this was thought to be a flaw in the models, but now there is good evidence that the physics of neutrinos was wrong and neutrinos change after being produced in the sun and before reaching the Earth. The neutrino deficit is now seen as a triumphant confirmation of the solar models
solar pulsations
The sun pulsates -- the pulsation periods are near 5 minutes. The pulsations cause the surface of the sun to be slightly corregated, and careful analysis of the corregations gives much information about the interior of the sun. Solar pulsations and solar models agree in two important areas:
• The boundary between the radiative and convective zones is at 0.7 solar radii.
• The abundance of helium is slightly higher inside the sun than at the surface.
sunspots
dark spots on the sun
magnetic fields
find in book
corona
layer above the suns visible surface
chromosphere
layer above the suns visible surface
solar wind
see book
evolution of a white dwarf
The envelope of the star gets used up, in part because the hydrogen in its envelope is fused to helium, in part because the star is losing its envelope through a stellar wind.
When the envelop is gone, the hydrogen-fusing shell is extinguished leaving the bare core of the star.
The radius of the core is about 0.01 solar radius and it contains perhaps 0.5 solar mass. It is a white dwarf!
Initially the core is hot but it cools off because it isn't producing new energy. It takes so long for a white dwarf to cool that no white dwarf has cooled below a temperature of about 2500 K.
evolution of a neutron star
• Stars with masses greater than about 8 solar masses evolve differently from low-mass stars. They evolve horizontally back and forth across the HR diagram as supergiants, finally becoming supernovae and ending up as neutron stars or black holes. Although rare, high-mass stars are important because they are the source of most of the heavy elements in the universe.
evolution of a black hole
• Stars with masses greater than about 8 solar masses evolve differently from low-mass stars. They evolve horizontally back and forth across the HR diagram as supergiants, finally becoming supernovae and ending up as neutron stars or black holes. Although rare, high-mass stars are important because they are the source of most of the heavy elements in the universe.
total disruption in a supernova explosion
• The iron core of an evolved high-mass star is supported by electron degeneracy pressure. When the core mass exceeds 1.4 solar masses, gravity overwhelms the pressure due to electron degeneracy and the core collapses.
o As the core collapses, electrons in the core join with protons to create neutrons ("electron capture"). Without electrons, there is no electron degeneracy pressure and the core goes into free-fall collpase.
o Collapse ends when pressure from neutron degeneracy becomes high enough to withstand gravity. The core becomes a neutron star. The collapse takes less than 1 second! 




o During the core collapse, the inner core collapses faster than the outer core. The inner core overshoots the equilibrium radius of a neutron star (it gets too small) and bounces back to larger radius.
o The bouncing inner core hits collapsing material from the outer core. The outer core is pushed outwards at high speed.
evolution of a typical star with a mass of 2.0 M
expanding envelope
of H and He
radius = 50 R
hin shell where
hydrogen fuses
to helium contracting

helium core
radius = 0.01 R
basic properties of white dwarfs and how they are formed
The core of the AGB star is left behind after the ejection of the outer layers. Without the overlaying layers, it rapidly cools, and no more fusion can take place. It is a dead star.
The radius of this dead star is about 0.01 solar radius and its mass is typically only about 0.6 solar masses. (It lost mass while a red giant and an AGB star, and when it ejected its planetary nebula.) It is a white dwarf.
The white dwarf starts off extremely hot (100,000 K or more), but it soon cools off. The cooling is rapid at first but the rate of cooling slows down and no white dwarf has yet cooled below a temperture of about 2500 K.
This kind of white dwarf is made of carbon and oxygen.
basic properties of neutron stars and how they are formed
• Stars with masses greater than about 8 solar masses evolve differently from low-mass stars. They evolve horizontally back and forth across the HR diagram as supergiants, finally becoming supernovae and ending up as neutron stars or black holes. Although rare, high-mass stars are important because they are the source of most of the heavy elements in the universe.
behavior of binary stars with mass transfer onto white dwarfs, neutron stars, and black holes
• A black hole can in principle have any mass. Neutron star, however, have masses less than 2.7 solar masses and white dwarfs have masses less than 1.4 solar masses
• If two stars in a binary star system are close together one can pull mass away from the other
supernovae: what stars look like before they erupt in supernova explosions, what causes the explosion, and their effects on the rest of the galaxy, especially on the chemical evolution of the galaxy.
• The iron core of an evolved high-mass star is supported by electron degeneracy pressure. When the core mass exceeds 1.4 solar masses, gravity overwhelms the pressure due to electron degeneracy and the core collapses.
o As the core collapses, electrons in the core join with protons to create neutrons ("electron capture"). Without electrons, there is no electron degeneracy pressure and the core goes into free-fall collpase.
o Collapse ends when pressure from neutron degeneracy becomes high enough to withstand gravity. The core becomes a neutron star. The collapse takes less than 1 second! 




o During the core collapse, the inner core collapses faster than the outer core. The inner core overshoots the equilibrium radius of a neutron star (it gets too small) and bounces back to larger radius.
o The bouncing inner core hits collapsing material from the outer core. The outer core is pushed outwards at high speed.
o Also, a huge number of neutrinos are created during the core collapse. The density in the core is so high that the neutrinos can interact with matter and the pressure from the neutrinos helps to eject the outer core. 




• The outer layers of the star are ejected at speeds of 10,000 to 30,000 km/sec (eg, at speeds up to 10% of the speed of light. In fact, even the best models for supernovae cannot fully explain why the outer layers are ejected, but we know that the layers are ejected because we see the ejected material (eg, supernova remnants).
• Most of the visible light from a supernova comes from the ejected material. The material takes months to fade away.
• Large bursts of neutrinos are given off by supernovae. The neutrinos have been detected by neutrino telescopes and are proof that core collapse causes the supernova.
the basics of special relativity
o The Special Theory of Relativity is based on three properties of light:
• The speed of light is 3 x 1010 cm/s.
• The speed of light is the same for everyone.
• Nothing can go faster than the speed of light.
describe how the solar system was formed
see book
how planets are formed
• A slowly-rotating cloud of gas and dust begins to collapse to form a star.
• The cloud's rotation speed increases as it collapses, flattening the cloud and forming a disk.
• A star forms at the center of the disk. The disk is now called a proto-planetary disk.
• Planets form in the proto-planetary disk.
o The remnant of the disk falls into the star or is blown off into space.
connections between planets, solar system and their formation
see book
general relativity
• Mass bends space-time in its vicinity; large masses bend space-time more than small masses.
• A particle moves in a straight line through this curved space time.
• We perceive this straight line motion in curved space-time as curved motion in ordinary space, and call it motion influenced by gravity
black holes
o Like the balls, the paths photons are straight lines in spacetime but curved in ordinary space and time. In other words, gravity bends light!

o Nothing, not even light, can escape a black hole.
o The black hole does not have a solid surface. The material of the original star is compressed to a point by the infinite gravity.
o The radius at which the gravity becomes infinite is called the radius of the black hole. The radius is 3 km for a black hole with a mass of 1 solar mass.
o A black hole can have any mass and radius, but its mass and radius are related by the equation
o R = 3 M
o where the mass is in solar masses and the radius in kilometers.
• As seen by a distance observer, time slows down near the edge of a black hole.