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Astronomy 1 > Stellar Physics > Flashcards

Stellar Physics Flashcards

1
Q

Method to find distance to nearer stars

A

Parallax

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2
Q

Luminosity of a star

A

It’s energy output per unit time
Measured in watts
If star radiates isotropically flux equal across sphere
L=4pir^2F

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3
Q

Visual magnitude

A

Magnitude in the visual part of the spectrum

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4
Q

If a star is hotter than its environment

A

It will cool down by re-radiating it’s energy

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5
Q

Rayleigh jean law for blackbodies

A

Good approximation at Lon wavelengths but radiance keeps increasing indefinitely at short wavelengths (UV catastrophe)

Failure of classical physics to explain thermal radiation

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6
Q

Wien’s law

A

Good approximation to the observed spectrum at short wavelength

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7
Q

Planck’s law

A

Assuming energy comes in discrete quanta
Fit data

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8
Q

What determines apparent colour of an object

A

Shape of spectrum and position of its peak

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9
Q

Atmospheric extinction

A

Light from sun obscured by atmosphere of Earth, which absorbs more than others.
Ignore in this course

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10
Q

Stefan Boltzmann law

A

Total luminosity per area is the spectral radiance integrated over solid angle and wavelength

j=L/A=sigma T^4

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11
Q

Total luminosity of a spherical star

A

Multiply Stefan Boltzmann law by surface area

L=4piR^2sigmaT^4

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12
Q

Stars have wide range of luminosity so helpful to use logs

A

Take logs of both side of luminosity equation
Plot graph of y=log10L against x=-log10T for some value r and get straight line

Lower radii lie on lines nearer bottom

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13
Q

HR diagram

A

80-90% of stars cluster in main sequence
Other branches of stars: white dwarfs, giants and supergiants
Temperature x axis decreases from left to right
Y axis is luminosity

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14
Q

Interpretation of HR diagram

A

Only certain combinations of L and T allowed
Most stars on main sequence

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15
Q

What tells us that stars move around HR diagram as they evolve

A

Clusters are stars at similar stages of their lives
Number of stars in each part proportional to duration of that stage of evolution

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16
Q

Main sequence

A

Most variation from top left to bottom right along a line of roughly constant radius
Top left blue stars hotter and more luminous

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17
Q

Two cut offs to the main sequence

A

At top: luminous stars blow material from their surface through radiation pressure naturally limiting their mass

At bottom: cool red stars not hot enough to begin nuclear reactions. Temperature in core too low

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18
Q

Estimate of time a star spends on main sequence

A

Lifetime=energy available/ luminosity

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19
Q

Most massive stars spend…

A

Least amount of time on main sequence

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20
Q

Giants and supergiants

A

Sit in top right of HR diagram
Large L but low T
Less populated so stars spend less time in this phase
Reach after main sequence

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21
Q

White dwarfs

A

Bottom left
Below main sequence so radius s smaller
None visible to naked eye
Are not powered by nuclear fusion

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22
Q

Photometric system

A

Divides spectrum into commonly used bands
Ultraviolet band centred 350nm, blue band 440nm, visible 550nm, red 600nm, near infrared 800nm

Filters placed over telescope to select a band

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23
Q

Colour index

A

Numerical difference in magnitudes between measurements made in two wavelength bands

Measurements made through two different filters eg B-V difference between magnitude in blue and visible band

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24
Q

The smaller the colour index (ie lower position on number scale that ranges from positive through zero into negative)

A

The more blue and hotter the star

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25
Q

Kirchoff’s first law

A

A hot and opaque solid, liquid or highly compressed gas emits a continuous black body spectrum with no spectral lines

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26
Q

Kirchoff’s second law

A

A hot, transparent gas illuminated by a continuum source produces a spectrum of bright emission lines

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27
Q

Kirchoff’s third law

A

If a continuous spectrum passes through a transparent gas at a lower temperature the cooler gas will absorb at characteristic wavelengths resulting in dark absorption lines

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28
Q

Harvard classification

A

Spectra of light from stars fell into natural categories based on strength of certain key line features

Every star has a letter that describes its colour, known as spectral class

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29
Q

Harvard classification scheme

A

Considers changes in other lines as well as hydrogen, gives sequence indicating source temperature

OBAFGKS

O hit and blue, m cool and red

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30
Q

O spectral class

A

Hottest blue stars
Few lines
Strong He II absorption lines

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31
Q

M spectral class

A

Coolest red stars
Spectra dominated by molecular absorption bands especially TiO
Strong metal absorption lines

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32
Q

Harvard subclasses

A

Each type is divided into 10 subclasses
These reflect gradual temperature change
eg for A0,…,A9
0 is hotter end
9is cooler end so O9 is next to B0

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33
Q

What does allocation of subtype depend on

A

Line strengths and ratios

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34
Q

How does the Harvard classification scheme not completely describe a star

A

Cannot distinguish between stars with the same temperature but different luminosities

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35
Q

Morgan keenan luminosity class

A

Established to add discrimination on the basis of luminosity

Ranges from I to VII

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36
Q

Class I

A

Hypergiant
Divided I-O
Through 1a, bright supergiant and 1b, dim supergiant

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37
Q

Classes II - V

A

Goes from bright giants down to main sequence dwarfs

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38
Q

Classes VI and VII

A

VI sub-dwarf
VII white dwarf

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39
Q

What are luminosity classes determined from

A

Mainly from observed width of spectral lines

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40
Q

Broadening of spectral lines

A

Several effects can cause broadening
High pressure and temperature causes atoms to collide more frequently which broadens spectral emission particularly in hot dense stars like white dwarfs

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41
Q

Mass luminosity relationship for main sequence stars

A

L/lsun ~ (M/m sun)^a

Value of a depends on fit used in data but approx 3<a<3.5
More massive, more luminous
Most massive in top left of HR

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42
Q

Explaining mass luminosity relationship

A

Massive stars have large gravitational compression of cores
For equilibrium, need high radiation pressure outwards
High thermal pressure provided by high temp in core
Nuclear reaction rate very sensitive to core temp so even slight change produces large change in luminosity

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43
Q

Implication of mass luminosity relationship

A

3-3.5 is big power
Great implications on how long star lives on main sequence
Massive stars have short lifetimes because they burn up fuel quicker

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44
Q

What does variable mean

A

Star’s flux changes over time
Observe by measuring changes in apparent magnitude

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45
Q

Real and apparent variation

A

Real: star itself changes
Apparent: eg something moves in front of star and blocks light partially or fully

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46
Q

Irregular and regular variation

A

Irregular: no particular pattern, can be sudden or random
Regular obviously opposite

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47
Q

Novae

A

Sometimes called cataclysmic variables
Flare in brightness irregularly
Luminosity can increase factor 1000 over period ~ a week
All novae exist in binary systems in which material transferred from one to another causing bright outburst

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48
Q

Grange point

A

Point in middle of binary system

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49
Q

Disk

A

Best way to distribute angular momentum

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50
Q

T Tauri stars

A

Class of irregular variables
Luminosity increases by factor 3 in few days
Very young powered by gravitational energy as they contract

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51
Q

Type II supernovae

A

End state of a very massive star
After nuclear fuel exhausted core collapses and outer layers blown off
Small dense neutron star remains surrounded by expanding spheres of circumstellar matter

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52
Q

Type 1a supernovae

A

Expect one type 1a SN every few decades

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53
Q

Types of supernovae

A

Classified according to observational features
1a all have nearly equal brightness - standard candles

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54
Q

Examples of regular variable stars

A

Cepheids

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55
Q

Cepheid variable stars

A

Very luminous giant or supergiant
Luminosity varies by factors up to 10
Variation repeats over periods between 1 and 100 days
Eg Polaris period ~4 days

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56
Q

Radial pulsation results in a regular pulsation of

A

Velocity of star’s surface
Effective temperature
Luminosity

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57
Q

Instability strip

A

Where cepheid variables sit in HR diagram
Lies at roughly right angles to main sequence towards giant branch
Stage on the way to being a giant

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58
Q

Cepheid period luminosity relation

A

Period of pulsation depends only on average luminosity of the star
Longer pulsation period, the more luminous the star

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59
Q

Two types of period magnitude relations

A

Type I - massive young cepheids: M=-1.8+2.4log10P

Type II - older smaller cepheids: same but 0.4 instead of 1.8

period in days

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60
Q

Cepheid variables as distance indicators

A

Observing cepheid, measure period of oscillation, find intrinsic luminosity from period luminosity relation

Measure flux, can use F=L/4pid^2 to calculate distance

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61
Q

4 main classes of binaries

A

Visual
Astrometric
Spectroscopic
Eclipsing

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62
Q

Visual binaries

A

With telescopes, possible to resolve two components

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63
Q

Astrometric binaries

A

Cannot resolve individual stars but where we see a periodic wobble of observed overall position

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64
Q

Spectroscopic binaries

A

Components are not resolvable but Doppler shits in spectral lines reveal there are two stars orbiting same centre of mass

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65
Q

Eclipsing binaries

A

Not resolvable but see periodic brightening and dimming

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66
Q

Binary system orbital analysis

A

For cases where one mass&raquo_space; other, deal with orbit as if larger mass is stationary
Masses of components can be more comparable. Newton and Kepler still apply

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67
Q

Binaries: if the masses are equal

A

The centre of mass is halfway between them

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68
Q

Binaries: if two masses are different

A

Centre of mass closer to heavier object
m1/m2=r2/r1=a2/a1

Where a is semi major axis

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69
Q

How to find semi major axis

A

If measure period, angular separation of stars and distance to binary is known

Allows to find total mass of system

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70
Q

Visual binaries finding mass

A

m1a1=m2a2
For visual, can measure a1 and a2 so can find m1/m2
Since we know total mass, can deduce individual masses

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71
Q

Redshift and blueshift - spectroscopic binaries

A

Star moves away, redshifted
Towards, blueshifted

Eg if A blue shifted and B red shifted then A towards and B away
No Doppler shift during tangential motion

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72
Q

Speed toward/away

A

Reduced by inclination of orbital plane
l is inclination angle v=vtrue sinl

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73
Q

Finding combined mass for spectroscopic binaries

A

v1/v2=a1/a2=m2/m1
Using v=row and w=2pi/T

v1+v2=2pia/T
Gives m1+m2=T/2piG(v1+v2)^3

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74
Q

Stars will eclipse each other only if

A

We are viewing the system near edge on (l~90 degrees)

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75
Q

Eclipsing binary: most light when

A

No overlap

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76
Q

Plot of magnitude against time for eclipsing binaries

A

Get dip at secondary minimum and bigger dip at primary minimum

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77
Q

Primary minimum

A

Hotter star moves behind the cooler

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78
Q

Secondary minimum

A

Cooler star behind the hotter

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79
Q

If T2>T1

A

Then F’>F1 so secondary minimum

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80
Q

T2<T1

A

F’<F1 so primary minimum

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81
Q

In eclipsing, spectroscopic binaries we can get

A

Orbital period from either light curve or spectroscopy
Speed of stars in orbit from spectroscopy which can be used to find size of orbit
Enough info to calculate mass of two stars

Light curve allows us to compute size of each star, by measuring time of transit and combining with speed measured in Doppler shift

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82
Q

Vogt-Russell theorem

A

Structure of a star, in hydrostatic and thermal equilibrium, with all energy derived from nuclear reactions, is uniquely determined by its mass and the distribution of chemical elements throughout its interior

83
Q

What do Kirchoff’s laws tell us?

A

the absorption lines indicate the presence of a cooler layer of diffuse gas, on top of a hot layer of denser gas

84
Q

Structure of sun

A

(From inner to outer)
Interior, photosphere, chromosphere, transition region, corona

85
Q

When we observe the sun, we see light from

A

Photosphere, chromosphere and corona

In descending order of contribution

86
Q

Photosphere

A

Sun has no solid surface but becomes opaque to visible light at the photosphere
T varies throughout
Most light comes from 5800K region where density is 1000th of air

87
Q

Sunspots

A

Photosphere marked by darker, cooler sunspots
Around 1500K cooler than rest of photosphere
Sunspots appear to move across suns disk, shows sun is rotating

88
Q

Rotation of sun

A

Ball of gas so different regions rotate at different rates

Equatorial region rotate in about 24 days
Polar regions take more than 30 days

89
Q

Solar flares

A

Extends hundreds of thousands of miles above surface
Originate in photosphere
Produce bursts of radiation across EM spectrum

90
Q

Chromosphere

A

Get little light from chromosphere
Seen as dim reddish pink glow, only really seen during eclipse usually light too weak to be seen against photosphere

91
Q

Corona

A

Very hot
Outer atmosphere extends as far as 3 solar radii
Only really seen from Earth during eclipse
Eventually leads to solar wind

92
Q

Spectral line strength depends on

A

Number of atoms present and temperature of gas

93
Q

Effects on lines due to bulk motion

A

Allow measurements of overall rotation and large scale expansion or contraction

94
Q

Effects on lines due to random thermal motions of atoms in atmosphere

A

Allow measurements of temperature and pressure

95
Q

Energy of nth level of hydrogen

A

E proportional to -1/n^2

  • because binding energy
96
Q

Transition between level m (higher) and level n

A

Photon would be emitted of energy corresponding to energy difference between states

deltaEmn prop. ((1/n^2-1/m^2)

97
Q

Calculating wavelength of emitted photon

A

Using rydberg equation

wavelength must be positive so use n<m

98
Q

Lyman series

A

Transitions from excited states to ground state

99
Q

Balmer series

A

Transitions from m>2 to n=2

100
Q

Emission

A

Electron drops to lower level
Photon emitted, of energy corresponding to e difference between transition levels

101
Q

Absorption

A

e jumps to higher level
Photon of correct energy needed for this to happen

102
Q

Strongest visible line from sun

A

Ha line

e going from n=2 to m=3

103
Q

Boltzmann factor

A

Probability of an electron occupying energy state E
- tells us states with E»kbT are very improbable

104
Q

If there are 8 possible states in n=2 shell

A

8 times the probability of finding an electron there

Can calculate ratio of population of atoms in two states

105
Q

Rydberg equation for other elements

A

Fails to take account of electron orbital screening in beaver elements
Single electrons around heavier nuclei, can multiply by Z^2

106
Q

Ionisation

A

Under extreme conditions
Lose electrons if absorb enough energy from photons
Can find ionisation energy by using Rydberg equation with m=infinity

Gives x=13.6/n^2 eV

107
Q

How do ions form

A

Absorption: of a photon with at least x joules of energy
Collision: with another particle like an electron (scattering)

108
Q

Ion formation by photon scattering

A

Average photon has E=kbT but some have higher energy
Only photons with energy > x can cause ionisation (same as higher chem)

109
Q

Ion formation by scattering

A

Collisions occur with other particles. Similar distribution of energies
Average E=3/2kbT

To find proportionof atoms ionised, find equilibrium in: photon +atom (reversible arrows)electron+ion

Equilibrium is Saha equation

110
Q

Saha equation

A

No of atoms/ no ions ~ 10^21T^3/2exp(-x\kbT) / no of electrons

111
Q

Rule of thumb for Saha equation

A

50% ionisation occurs when kbT~x/18

112
Q

If a gas gets too hot

A

All atoms may already by ionised

May not be any low level electrons available to absorb
May not be any high level electrons able to emit

113
Q

If a gas gets too cool

A

Electrons may be in too low an energy state for a particular line
(Balmer needs n=2 for example)

114
Q

Effects of temperature on helium lines

A

Little seen from sun in absorption although they are observed faintly in upper chromosphere during eclipses
Helium ions can be excited to a state which gives visual absorption lines if temp right

115
Q

Metal lines (metal in Astronomy is anything after helium)

A

Very rare
Only dominate at low temperature where H and He frozen out
Relative abundances of metals in all stars is fairly similar
Abundance of metals relative to hydrogen is very different in some stars

116
Q

Population I stars

A

Ratio of metals to H and He is similar to Sun

Young stars, generally made from material ejected from older stars
Formed late in evolution of galaxy and are found predominantly in galactic disk

117
Q

Population II stars

A

Ratio of metals to H and He is 100 times less than that found in sun

‘Metal poor’
Old stars formed before galaxy was a disk and are found predominantly in galactic halo

118
Q

Population III stars

A

Possible massive stars formed just after Big Bang
Obviously no metals around then

119
Q

Molecular bands

A

Molecules can form in outer atmosphere of cool stars

Typical binding energy of molecule is 4-6eV

This means lines will fade if T>~5000K

We observe TiO, ZrO, CN and sometimes H2O

120
Q

Solar energy requirements

A

Know mass,luminosity of sun
Age of solar system is ~4.6 billion years

Energy required to keep sun shining is E~tau L where tau is solar lifetime

121
Q

Efficiency of sun’s energy

A

e=E/M> or = tauL/M

Units if j/kg

122
Q

Candidates for source of suns energy

A

Chemical energy
Gravitational energy
Relativistic energy
Nuclear fusion

123
Q

Chemical energy

A

Chemical reactions involve rearrangement of electrons
Units is electron volts
Comparing to mass per atom
Efficiency is around 100million j/kg so if chemical reactions were source of power, would only last Me/l =50000 year

124
Q

Gravitational energy

A

Using gravitational energy for a proton coming from infinity to r E=GMm/r so e=Gm/R ~10^11 j/kg

Still not enough to power sun but considerable amount of energy

125
Q

Relativistic energy

A

Using eff=mc^2/m gives 9x10^16 j/kg
tau=mc^2/L = 3x10^13 years

Easily long enough to keep sun bright
However, this conversion requires equal amount of matter and antimatter

126
Q

Nuclear fusion

A

Very high temp and density, several light nuclei can fuse to form a single nucleus
Newly formed nucleus will have slightly lower mass
Delta m is released as fusion energy

127
Q

Nuclear fusion in the sun

A

Four H nuclei (protons) fuse together to form a Helium-4 nucleus (alpha particle)

Energy released in the form of radiation (gamma rays) and neutrinos

128
Q

Energy budget

A

DeltaE=deltam c^2

Calculating delta e is energy available in one H-He fusion reaction
Eff=deltaE/4mproton, tau=Msuneff/L gives 400 billion years

This is big but plausible if we assume not all of sun’s mass undergoes fusion

129
Q

Interstellar gas cloud

A

Galaxy full of H and He from Big Bang and dust and gas from previous generations of stars

If not completely uniform, gravity will cause clumps to collapse
As it collapses, gas heats up(potential energy converted to heat energy)

130
Q

If cloud is initially rotating even slightly

A

Rotation speed Increases as material collapses inwards

131
Q

Spinning material

A

Flattens into protoplanetary disc with dense central part forming a protostar

132
Q

Surrounding material accretes onto the protostar and as temp rises,

A

Gas becomes ionised

133
Q

Eventually, if enough material accreted, temp becomes high enough to

A

Initiate nuclear fusion
Radiation from the hot starpishes material out until equilibrium is reached

134
Q

How much energy would it take to unbind cloud

A

Approximate two half spheres each M/2, distance R apart
Use equation for potential
Assume uniform density
Ereleased=Einit-Efin
If very large at start assume Einit=~0 so energy available is Gm^2/R

135
Q

Why do we not see many protostars

A

Lifetime if powered by left over heat from contraction is Kelvin helmholtz timescale

About 10^7 years for sun which is quite short

136
Q

Brown dwarf

A

Pressure and temperature never high enough for nuclear fusion
Will shine for its kelvin helmholtz timescale

137
Q

Hayashi tracks

A

Evolutionary path for protostar moves it around on HR diagram
For low mass stars, vertical motion downward until proto star hits main sequence
unstable during this phase

138
Q

Henyey tracks

A

For higher mass stars
Continues to heat up, gas in cloud fully ionised
No atoms to provide atomic absorption lines so gas transparent in outer regions
Protostar slowly contracts until onset of nuclear reaction in core

Nearly horizontal track on HR (slight luminosity increase with large temperature increase)

139
Q

Distinctive characteristics of pre main sequence stars

A

Unstable luminosity
Eject lots of gas
Surrounded by warm clouds

140
Q

Bipolar outflows

A

Young stars
Jets of material in opposite directions extending over a distance ~light year

Indicate new star is gaining material from a surrounding accretion disk

141
Q

Emission nebulae

A

Young stars often surrounded by emission nebulae
UV light from hot star sweeps out a gravity and ionises surrounding hydrogen
Recombination produces Ha

142
Q

Virial theorem

A

Total kinetic energy of a stable, self gravitating mass distribution is negative one half of the total gravitational potential energy

143
Q

What is virial theorem used for

A

Helps calculate the conditions that must exist for cloud collapse

144
Q

Derivation of jeans mass

A

Number of particles with mass mp in cloud is N=Mc/mp
Potential energy of spherical cloud with uniform density
Gas so ask of each particle is 3/2kbT
Total Ekby multiplying by N
Use viral therorem
Cloud collapse if unbalanced so use <
Use initial density and sub in to find Mc

145
Q

Jeans mass

A

Critical mass

A cloud with Mc>Mj will collapse under gravity

146
Q

Cloud collapse temperature

A

Use virial theorem
3/2kbTM/mp=1/2 x 3/5 GM^2/R

147
Q

What provides the energy needed to stabilise the cloud against further contraction

A

Temperature released by fusion

148
Q

p-p chain reaction

A
  1. h1+H1 —> positron+neutrino+ H2 (positron denoted by beta)
  2. H1 + H2 —> He3 + gamma
  3. He3+He3 —> He4+2H1
149
Q

Net reaction for p-p chain reaction

A

4H1 —> He4 + 2 positrons + 2 neutrinos + 2 gamma

150
Q

What is needed for p-p chain reaction

A

Carbon
Acts as a catalyst but not used up

151
Q

For CNO cycle to work at all

A

Carbon nitrogen and oxygen must be present which depends on the star type

152
Q

When does pp chain dominate

A

Lower masses and temperature

153
Q

When does CNO cycle dominate

A

Higher temperatures

154
Q

Why is CNO more important for heavier stars

A

Energy production rate varies strongly with CNO cycle so is more important for heavier stars which have higher interior temperatures

155
Q

Are we justified in modelling stars as gases

A

Yes
At high temperature hydrogen is ionised into protons and electrons (TINY compared to H)
In photosphere density of H tiny and mostly unionised
Density of neutral H much lower than maximally packed H, behaves like an ideal gas

156
Q

Hydrostatic equilibrium

A

Inside a stable star, inward pull of gravity is balanced by the thermal pressure due to heat

157
Q

Hydrostatic equilibrium equation derivation

A

In equilibrium, forces on each side balanced
Fpressure=Fgrav

158
Q

Solving hydrostatic equilibrium equation

A

Differential equation
Solve to find pressure as a function of radius
Assume uniform density
Need boundary conditions: P(R)=0, at centre, r=0 let pressure be P0
Solve by integrating both sides

159
Q

What would make a more accurate estimate of hydrostatic equilibrium

A

Density increase towards centre

160
Q

Central temp in sun

A

Use ideal gas law
Rewrite in terms of P and R=Nakb
Simplify using Nana=M and divide by u (mean molecular weight)

161
Q

More realistic models of stellar structure

A

Allow for temperature, pressure and density gradients inside star
Convection can occur instead of simple radiation

162
Q

What transpires heat energy from core to surface in sun like stars

A

Both convection and radiation

163
Q

Heat transfer in stars <0.5 solar masses

A

Convection

164
Q

0.5-1.5 solar masses heat transfer

A

Radiation then convection

165
Q

Heat transfer >1.5 solar masses

A

Convection then radiation

166
Q

End states of stellar evolution

A

White dwarfs, neutron stars and black holes

167
Q

When fuel runs out

A

Star will resume contracting since there is no longer a thermal pressure to balance gravity

168
Q

Low mass stars

A

Stars dues elements up to carbon
After, because stellar mas low, core pressure due to self gravity is not enough to cause temp rise high enough for heavier elements

Carbon core collapses to form white dwarf, outer layers expelled to form planetary nebula

169
Q

High mass stars

A

Continue fusion up to nickel and iron
Iron is most tightly bound atomic nucleus, no more energy available from fusion
Star collapses but much more massive, more gravitational energy can be released, more violent event - supernova

170
Q

ZAMS

A

zero age main sequence mass

171
Q

Supernovae are classified into different types depending on

A

Whether they show hydrogen in their spectrum

Type I have almost no H, type II do show presence of H

172
Q

Type Ia supernovae are believed to be the result of

A

Accreting white dwarfs being pushed over the mass limit and undergoing further nuclear fusion in a runaway reaction that destroys the star

173
Q

Types Ib, Ic and II result from

A

Collapse of high mass stars straight to a neutron star or black hole

174
Q

SN with no H but has Si

A

Type Ia

175
Q

SN with no H, no Si but does have He

A

Type Ib

176
Q

SN with no H, Si or He

A

Ic

177
Q

SN with H

A

Type II
(First thing to lo9 for are Balmer lines

178
Q

What are type Ia SN characterised by

A

Silicon

179
Q

Where are types Ib and Ic found

A

Only in spiral galaxies where there was recent star formation (implies they are due to massive short lived stars)

180
Q

Onion like shell structure with fusion by products

A

Final stage is silicon burning, generates host of nuclei centred around 56Fe minimum of the 26 binding energy curve

Each stage less energy efficient (take shorter and shorter) eg H burning 107 years, Si ~days for star of 20 solar masses

181
Q

Nuclear fuel exhausted in core, radiated energy can no longer be replaced so

A

Pressure falls, hydrostatic equilibrium lost, core rapidly contracts

Releases g potential energy, T rises allowing endothermic reactions

182
Q

Photodisintegration

A

Earlier exoteric fusion reactions are reversed

183
Q

Neutron capture

A

Endothermic reactions create elements heavier than Fe by neutron capture
Can only get elements heavier than Fe because of supernovae

184
Q

Electron capture

A

Endothermic reactions absorb gravitational energy released by initial collapse, further contraction, further temp rise
Core can get hot enough for inverse beta decay
Causes extreme pressure drop in core
Core separates from outer envelope and goes into free fall

185
Q

Collapse

A

Closely packed neutrons
Core stiffens due to degeneracy pressure and abruptly halts collapse
Sends shock wave back up through in falling material
Shock wave initiates further endothermic reactions and loses energy

186
Q

Remant

A

Bouncing material and enormous neutrino flux blows outer layers off star into space, forming supernovae remnant

187
Q

Why do SN remnants expand so quickly

A

Momentum transferred from lower layers to upper
Shock wave carries energy and momentum, encountering material of ever decreasing density
P=mv roughly constant p and smaller m so v big

188
Q

Luminosity- energy source is

A

the release of gravitational potential energy by the contraction of the core

189
Q

What do supernovae produce

A

Cosmic rays (extremely high energy particles, mostly protons)

190
Q

Supernovae distribute

A

Heavier elements into space which are incorporated into new stars and planets

191
Q

Pauli exclusion principle

A

Allows at most one fermion to occupy a given quantum state since no two fermions can have the same set of quantum numbers

192
Q

Degenerate

A

At zero temperature all of lower states and none of higher states occupied

193
Q

Heisenberg uncertainty principle

A

Particle confined to smaller and smaller volume in space will have a larger uncertainty in momentum

194
Q

Degeneracy pressure

A

In stellar remnant, central pressure must be balanced by the cold pressure

195
Q

Simplified model

A

Cubic meter of particles
Number density n so contains n particles
Pauli - different particles

196
Q

Uncertainty in position of particles in simplified model cannot be

A

Larger than their physical separation

197
Q

Spacing between each particle is delta x so volume

A

Delta x cubed

198
Q

Delta x =

A

1/n1/3 so n=1/V

199
Q

By HUP each particle has momentum in x direction of

A

Px ~ h bar/ delta x

200
Q

Average kinetic energy of non relativistic particles

A

E=1/2mv^2 (sub in momentum)

*p^2 = 3px^2

201
Q

Average quantum energy of electron

A

Sub in equation for momentum into Ek

202
Q

Since me«mp

A

Ed»
Electron pressure dominates over proton pressure for white dwarfs

203
Q

To find degeneracy pressure

A

Consider work done by changing volume of box by dx

Astronomy 1 (6 decks)

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