u. klein; j. kerp- physics of interstellar medium
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Physics of the interstellar medium
Prof. U. Klein
P.D. J. Kerp
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timeline
Today: Introduction (JK)
28. April 2009: line radiation, fundamentals, radiationtransfer (JK)
5. May 2009: continuum radiation processes (UK)
12. May 2009: heating and cooling (JK)
19. May 2009: dust (JK)
(UK)
Written exam: 21. July 2009 16:00 Hrsaal 0.03
Second date: 08. September 2009 10:00 Hrsaal 0.03
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Further reading
James Lequeux The interstellar medium Springer Verlag2005
Ewine van Dishoeck Master lecture Leiden University
2006 http://www.strw.leidenuniv.nl/~dave/ISM
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The interstellar medium:line radiation processes, physicalconditions of the gas
J. Kerp
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Outline
The interstellar medium in general
History of the discovery
Composition
Phases of the ISM
Transitions Basic physical quantities
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The interstellar medium in general
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The interstellar medium in general
The radiation of a galaxy isdominated by the stellar light
The mass of a galaxy isdominated by the stellar massdistribution
Studies of the stellar populationsof galaxies are essential todisentangle the formation history
of a galaxy
The radiation of the ISM is faintand the bulk of photons areemitted in the radio and infraredwavelength regime
The ISM hosts only a fewpercent of the total mass of agalaxy
The composition of the ISM iscontinuously modified by thestellar evolution
The distribution and chemicalcomposition of the ISM iscontinuously modified by
accretion of matter by the galaxy Density and temperature
variations within the ISMtrigger the star formation rateof a galaxy
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Galaxy: M82
www.wikipedia.org
Visible light/HAST
Starformation leadsto transfer of matter
outside the stellardistribution
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Galaxy: M82
www.wikipedia.org
Blue: X-ray/Chandra
Red: Infrared/Spitzer
Visible light/HST
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Dark clouds
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Dark clouds: Horsehead Nebula
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Dark clouds: Horsehead Nebula
Compiegne et al. 2008, A&A 491, 797
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Dark clouds: Horsehead Nebula
Compiegne et al. 2008, A&A 491, 797
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Dark clouds: Horsehead Nebula
Compiegne et al. 2008, A&A 491, 797
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The ISM in general
van Dishoeck ISM lecture
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The discovery of the interstellarmedium
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Stellar distribution
First attempts to determine the distribution of the stars werde made beThomas Wright (1750) and Immanuel Kant (1755). Systematic starcounts were performed by Wilhelm Herschel (1738-1822). He producedthe first map of the stellar density distribution. According to his star
counts, the Sun is thought to be in the very center of the Milky WayGalaxy.
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Stellar distribution
Astronomie II Gondolatsch, Groschopf, Zimmermann, Klett Studienbcher
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Stellar distribution
We assume,
that the stars fill entirely homogeneously the space.
All stars have the same luminosity.
The apparent magnitude of a star mis a function of r-2
. Accordingly, stars with an apparent magnitude in excess of mfollows
the relation
log r= 0,2 m + const.
The number of stars in space increases proportional to r3
Accordingly we find
log N(m) = 0,6 m+ const
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Stellar distribution
Neuer Kosmos Unsld, Baschek, Springer Verlag
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Stellar distribution
0,06
10
350
3000
0,25
50
6000
200.000
6,0
11,0
16,0
21,0
Polar region
[N/square degees]
Plane
[N/square degrees]
Apparentmagnitude
[mag]
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Stellar distribution
Considering the vertical distribution of the stars, we can approximatethis by a hydrostatic density distribution.
According to this approach, we can characterise the distribution by asingle quantity, the scale height . z denotes the vertical distancefrom the galactic plane and 0 is the mid-plane stellar density.
=
z
ez 0)(
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Z-distribution of the stellar populations
Astronomie II Gondolatsch, Groschopf, Zimmermann, Klett Studienbcher
Scale-heights for different types (Objektgruppe) of stars
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Discovery of the ISM
~ W. Herschel compiles thefirst catalog of "nebulae"
Until ~1900 absorption lines
have been discovered, but it isunclear whether they are ofstellar (circum-stellar) orinterstellar origin
1919 Barnard compiled thefirst catalog of "Dark Clouds"
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Discovery of the ISM
1933 Plasket & Pearce found a correlation between theCaII absorption line strength and the stellar distance.
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Discovery of the ISM
~1937 the first interstellar molecules CH, CH+ and CN werediscovered
1945 van der Hulst predicted the detect ability of the HI 21-
cm line 1949 discovery of interstellar magnetic field by polarization
measurements
1950`s maps of the Milky Way in HI (10% of the stellarmass is in HI)
LAB Kalberla et al. 2005
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Discovery of the ISM
1945 van der Hulst predicted that the hyperfinetransition ofthe radiation emerging from the hydrogen ground level willbe detectable within the Milky Way
1951 Ewen & Purcell and Oort & Muller detectedidenpendently the HI 21-cm line of neutral hydrogen.
This was the starting point to explore the Milky Way Galaxy
in HI 21-cm emission. The HI disk is much larger than the stellar disk
The total gas mass of the HI disk is about 10% of the stellar mass
The average volume density of atoms is 1 cm-3
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Properties of the HI 21-cm line
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Detection of the 21-cm line
http://www.astro.uni-bonn.de
accuracysurface2cm
dish100mawith9dish25mawith53
D
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Distribution of the line intensity
http://www.astro.rug.nl
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Properties of the HI 21-cm line
Natural line width is 10-16 km s-1, accordingly the line isideal suited to trace turbulence and Doppler motions.
The thermal line width is ~ 0.09 T, which corresponds
to 2 km s-1 at T= 100 K Doppler shift
=
=
c
c1
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HI in the Milky Way
HI 21-cm rotation curve of the Milky Way galaxy. Shown is the velocity measured relative to the
local standard of rest versus the galactic longitude (Burton 2001)
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HI in the Milky Way
Assuming that the gasencircles the center ofthe Milky Way on
concentric orbits, it isfeasible to determine isdistance (tangentialpoint method).
The method isrestricted to the innergalaxy. However, itprobes a large fraction
of the baryonic mass.
van Dishoek lecture 8
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HI in the Milky Way
van Dishoek lecture 8
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HI in the Milky Way
The tangent point method is only applicable
for R < RSonne.
The outer galaxy can be explored by the HI21-cm line only via continuing the rotation
curve. The derived dynamical distances are
only estimates with partly highuncertainties.
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HI in the Milky Way
Kalberla, 2003, ApJ 588, 805
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HI in the Milky Way
Velocity
Expected line shape
Observed line
van Dishoek lecture 8
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Discovery of the ISM
http://www.astron.nl
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Discovery of the ISM
1960 discovery of the soft X-ray background
http://www.mpe.mpg.del
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Discovery of the ISM
1963 the first interstellarmaser had be discovered(OH)
1968 NH3, the"thermometer" in theUniverse was observed forthe first time
1970 12CO(10), thesecond most abundantmolecule in the Universe
was discovered
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Discovery of the ISM
1970`s infrared
astronomy opens the
window to the most
abundant molecule H2
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H2: necessity for tracers
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Discovery of the ISM
~1990 Sub-millimeterastronomy opened thewindow to molecular clouds
and star forming regions 1990 COBE studied the
distribution of the dominantcooling line of the ISM CII.
1995 allowed detailedspectroscopic studies of the
dust, the vibrationally
excited H2 emission lineand the infrared dark clouds
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Discovery of the ISM
S
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Discovery of the ISM
Infrared Dark Clouds
Di f h ISM
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Discovery of the ISM
Di f th ISM
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Discovery of the ISM
1998-2006 SWAS studied thedistribution of H2O, O2, CI upto 500 GHz
2000-today FUSEobservation of H2 inabsorption againstbackground continuum
sources, observation of thevertical structure of highlyionized gas like OVI and NV
2003-today Spitzer studiesthe ISM with high angularresolution
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Basic properties of the ISM
Basic properties of the ISM
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Basic properties of the ISM
Confined to the Galactic Plane (much flatter than acompact disk!)
The interstellar medium consists mainly of hydrogen and
helium. All elements heavier than hydrogen are denoted as metals
Temperature range 4 K < T < 106K
The temperature is used as a measure for the physical conditions of
the interstellar gas. The phases of the interstellar medium arecharacterized by the average gas temperature.
Densities 10-4 cm-3 < n < 107 cm-3
Far from thermal equilibrium!Not easy to calculate!
Abundances of elements
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Abundances of elements
4.310-5Fe2.210
-6
Ca
1.710-5S
4.310-5Si
3.110-6Al4.210
-5
Mg
2.110-6Na
4.510-4O
6.310-5N2.510
-4
C
0.075He
1.00H
AbundanceElement
Abundances of the ISM
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Abundances of the ISM
"Metalls"He about 10%
C,N and O
Si, Ca and Fe bound in dust grains
Grains
About 1% of the ISM mass
Photons
Magnetic fields
Cosmic rays
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Phases of the interstellar matter
Classification of the ISM
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Classification of the ISM
Chemical composition of the ISMcomparable to the elements abundance of
the Solar System The state of Hydrogen determines the state
of the ISM
Molecular region H2Neutral region HI
Ionized region H+
Molecular regions
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Molecular regions
Molecular clouds aredefined by the presenceof molecular hydrogen.
We differentiate between: Diffuse molecular clouds
T 40 ... 80 K
n 100 cm-3
Dark clouds
T 10 ... 50 K
n 104 ... 106 cm-3
Dust is an importantconstituent of molecularclouds.
Molecular regions (sub-mm range)
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Molecular regions (sub-mm range)
van Dishoeck ISM lecture 1
Molecular regions (sub-mm range)
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Molecular regions (sub-mm range)
van Dishoeck ISM lecture 1
Molecular regions (UV-range)
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Molecular regions (UV range)
van Dishoeck ISM lecture 1
Neutral gas regions (radio-range)
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Neutral gas regions (radio range)
Dusty cirrus clouds
T 80 K
n 1 cm-3
Warm neutral gas
T 6000 K
n 0.05 ... 0.2 cm-3
Neutral gas regions (infrared range)
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Neutral gas regions (infrared range)
Neutral gas regions (infrared range)
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Neutral gas regions (infrared range)
van Dishoeck ISM lecture 1
Ionized gas regions
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o ed gas eg o s
HII regions are envelopesof early-type stars
T
104
Kn 0.1 ... 104 cm-3
Ionized gas regions
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g g
van Dishoeck ISM lecture 1
Ionized gas regions
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g g
van Dishoeck ISM lecture 1
Ionized gas regions
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g g
Coronal Gas T 106 K
n 0.005 cm-3
Ionized gas regions
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g g
van Dishoeck ISM lecture 1
Phase transitions: Orion nebula
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Phase transitions
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van Dishoeck ISM lecture 1
Phase transitions: M82
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Blue: X-ray/Chandra
Red: Infrared/SpitzerVisible light/HST
Phase transitions
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Physical quantities
Atoms and Ions
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The dominant part of baryonic matter in theuniverse is in neutral or partly ionized state (Warm
Hot Intergalactic Medium, 80% of all baryons)
Line emission and absorption of atoms and ionsare the dominant sources of photon emission and
absorption
Basic knowledge of atomic physics is necessary tounderstand the processes in the universe
Fraunhofer lines
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Astronomie I, Gondolatsch, Groschopf & Zimmermann, KlettStudienbcher
Fraunhofer lines of the Sun, from 390 nm to 690 nm
Fraunhofer lines
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Fraunhoferlines can be observed in laboratory experiments. Acontinuous light source is observed through a gas. Using a
spectrometer, absorption lines are observed, which are unique for
individual chemical elements.
If the continous light source is switched off, the absorption lines changeinto emission lines in the spectrum.
Fraunhofer lines originate via resonance absorption. The gas
attenuates the continous light emission only at those wavelengths
(energies) which corresponds exactly to the spacing of the energylevels.
Linien radiation
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Top: continous spectrumMiddle: emission spectrum of the gas
Bottom: absorption spectrum of the gas
http://de.wikipedia.org
Line radiation processes
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http://www.homepages.ucl.ac.uk/~ucapphj/EinsteinAandB.jpg
Line radiation: hydrogen
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Line radiation: hydrogen
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Line radiation: hydrogen
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http://pulsar.sternwarte.uni-erlangen.de/wilms/teach
Line radiation: line width
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The natural line with is determined by the time-energy uncertainty dueto the Heisenberg uncertainty principle
This relates the life time of an excited state with the energy uncertaintyof the transition. The natural line width can be describe by a Lorentzian
profile
determines the FWHM and the amplitude 1/
tE 2
h
( )( )22
,
+
=x
xL
Linien radiation: line broadening
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Astronomie I, Gondolatsch, Groschopf & Zimmermann, KlettStudienbcher
Radiation from a background sourceis absorbed by the gas cloud G. Theabsorbend photon energy is emittedisotropically. The total amount of
absorbed energy is relased after theabsorption processes but withoutany preferred direction.
Observer
Line radiation: line broadening
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Inelastic scattering
Maxwell-Boltzmanndistribution
Line radiation: line broadening
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An absorption line is produced wenna atom absorbs an amount of energy
E=hc/ at a wavelength .
The electron which is bound to the
electronic state Em is excited after the
absorption to the higher energy state
Ek.
ch=
mk EE
Equivalent width
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Astronomie I, Gondolatsch, Groschopf & Zimmermann, KlettStudienbcher
Definition of the full-width halfmaximum equivalent width .IK denotes the intensity of the
continuum at the wavelenghtof the absorption line. I0 is theintensity minimum and I is theresidual intensity at the
wavelength . A is thecorresponding equivalentwidth, which has the samearea as the the absorption line.
Full-width half maxium
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The advantage of use the full-width half maximum definition is,that we apply the assumption,that the absorbing/emitting gas
atoms are following a Maxwell-Boltzmann distribution.
According to this assumption wefind
( )
=
=
35482.22ln8
2
1)(
2
20
2
FWHM
exf
xx
Maxwell-Boltzmann FWHM
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( )
02
02
2
2
0
2
2
2ln8
2)(
2
)(
20
20
2
fmckTf
fmc
kT
dfekTf
mcdffP
de
kT
mdP
FWHM
f
kTf
ffmc
kT
m
=
=
=
=
Line radiation: line broadening
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We apply the non-relativisticDoppler formula
For a gas with the moleculare mass
M* we can calculate the average
velocity dispersion of the gas atoms
Applying the Doppler shift formula,
we can calculate the average line
width due to thermal motion.
A typical value for the interstellar
medium is
nm0026,0
kmol47,9kg
K106kmolKJ10314,82
sm103
m10491,5
R2
c
R2
c
1-
31-1-3
1-8
7
*
*
=
=
=
=
=
D
D
D
r
M
T
M
T
Line shape
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Astronomie I, Gondolatsch, Groschopf & Zimmermann, KlettStudienbcher
Doppler core
Damping wings
Damping wings
Line radiation: line broadening
A h i th i h b f 2
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As shown in the viewgraph before, wehave two processes which lead to the
broadening of an absorption/emission
line.
Thermal broadening (temperature T)
Pressure broadening (temperature T and
volume density n)
Thermal and pressure broadeing are
different in their physical behaviour.
Thermal broadening implies a folding of
the Lorentzian profile (natural line width)
with the Maxwell-Boltzmann velocity
distribution.
Pressure broadening leads to a shift of
the atomic energy levels due to the
interaction process. In gerenal, the
interaction with ambient gas atoms is onshorter time scale than the spontaneous
emission. Here, the volume density nis
of prime importance, Tis of second
interest.
( )
( ) 220
1
20
DI
eI D
+
Evaluation the strength of an absorption line
A photon of the wavelenght will be absorbed by a ion which has the ionzation degree i
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A photon of the wavelenght will be absorbed by a ion which has the ionzation degree iwill move the excited electron from orbit m (Ei,m) to orbit k (Ei,k).
The strenght of the absorption line is proportional to the path length through the
absorbing medium H.
The strenght of the absorption line is also proportional to the number of ions in the same
energetic state Ei,m Finally, the strength of absorption is a function of the propability that the absorption
events occurs quantum mechanically. This propability is describe by the oscillator
stength fm,k
HNfA mikm ,,
Evaluation the strength of an absorption line
1 Determination of A
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1. Determination of A.2. The gas atom, its ionization state i, energy levels Ei,mand Ei,kand fm,k
are known, accordingly A
yields directly Ni,mH
3. Using the Boltzmann equation one can calculate the population of
the Ions at a certain energy level (Ei,m
)
= T
EE
i
mi
imi
e
N
N k
1,
,
1,,
Evaluation the strength of an absorption line
We know from quantum mechanics that most of the energy levels
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We know from quantum mechanics, that most of the energy levelsshow up with a degeneracy in the orbital and magnetic quantum
number. These hidden energy levels are well known as fine-
structure levels. According to this degeneracy, we have to apply som
weightening to the number of energy levels.
= T
EE
i
mi
i
miimi
e
g
g
N
N k
1,
,
1,
,1,,
Fine-structur
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Evaluation the strength of an absorption line
Finally we need the number of elements which are within the excited
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Finally, we need the number of elements which are within the excitedstate Ei,m
The population of elements in the state is a function of the
temperature. It follows the Boltzmann equation but deals with the
ions. The equation is the Saha equation
=
TEE
eee
Tm
g
g
N
NN k3
1,0
1,1
1,0
1,11,01,1
h
k22
Evaluation the strength of an absorption line
Using the Saha and the Boltzmann equation we can calculate
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Using the Saha and the Boltzmann equation, we can calculateTemperature
Elektron density
Element abundance
Elektron pressure Accordingly, we can determine the physical state of a gas!
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