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Ch 8: Stars & the H-R Diagram
Nick Devereux 2006Revised 9/12/2012
Spectroscopy
Stellar Spectra
Emission & Absorption Lines
The Sun’s Spectrum
The Spectrum of Hydrogen
Electronic Transitions in the Hydrogen Atom
Energy Level Diagram
Two important facts about Hydrogen
1. The ionization potential = 13.6 eV2. The wavelength of the H emission line is 6563Å
where 1Å = 10-10 m
You can figure out everything about Hydrogen from these twofacts and knowing that the energy difference between two electronic states, E, is proportional to
E 1/n2
Where n is the principal quantum number
For Example
The Habsorption line results from the electron jumping fromthe n=2 to n=3 level. We can use this fact (#2) to calculatethe constant of proportionality, R
E = hc constant [ 1/ n1
2 – 1/n22 ]
So that, 1R [ 1/ n1
2 – 1/n22 ]
Substitute n1= 2 and n2 = 3 and = 6563Å
to yield R = 1.097 x 10-3
which is known as the Rydberg Constant
Now you can calculate the wavelength for all other
electronic transitions
Since, 11.097 x 10-3 [ 1/ n1
2 – 1/n22 ]
But, remember, that this equation yields in Å
Question: what is the wavelength of the Lyman transition ?
Question: what is the wavelength that corresponds to the ionization potential of hydrogen ?
Ionic SpectraThe wavelengths of H like ions, such as He II, Li III, O VIII,
can be estimated using the Bohr model of the H atom;
E = - U/2
E Z2 R[ 1/ n12 – 1/n2
2 ]
So, the energy levels scale by Z2 compared to H.
e.g. He has Z=2 , Z2 = 4, so the Lyman transition in He I
may be calculated using;
1Z21.097 x 10-3 [ 1/ n12 – 1/n2
2 ]
yielding = 304 Å which occurs in the far UV part of the spectrum.
Back to the Sun
The Balmer Discontinuity
The Balmer discontinuity is the break in the spectrum at 3646 Å due to the ionization of Hydrogen from the n=2state;
11.097 x 10-3 [ 1/ n12 – 1/n2
2 ]
Substitute n1 = 2, n2 = ∞, to yield = 3646 Å
Summary
Stars are made mostly of Hydrogen. So, the stellar continuumexhibits Hydrogen absorption lines which modifies it from aperfect Planck function.
The absorption line strengths depend on two things;
How many of the H atoms are in each of the excitation levels,n=1, n=2, etc, described by the Boltzmann equation
and secondly,
how many of the H atoms are completely ionized, for if theH atoms are ionized they can not produce any absorption lines.This is described by the Saha equation.
The Spectral Classification of Stars
The relative strength of the absorption lines depends on thetemperature of the star. If the star is too hot, all the H is ionizedand there are no absorption lines.
If the star is too cool, all the H atoms will be in the ground state, for which there are no transitions in the optical part of the spectrum.
A star of “medium” temperature will have lots of H absorption lines.
The range of H absorption lines strengths described above definesanother way to measure stellar temperature.
The Strengths of Absorption Lines
Boltzmann & Saha Equations
The strength of the absorption lines can be calculated fairly easily for H and with some difficulty for other species.
The calculation is a two step process.
Step 1 utilizes the Boltzmann equation to determine the relative number of atoms in the various energy levels with principle quantum level n =1, n=2, etc.
Step 2 utilizes the Saha equation to determine the relative number of atoms that are ionized.
Boltmann EquationN2 = g2 e –h/kT
N1 g1
Where:
N1 = number of atoms in n = 1 level
N2 = number of atoms in n = 2 level
g1 = 2 (statistical weight = degeneracy = 2n2)
g2 = 8
T = temperature
= frequency
Application of the Boltzmann Equation
For the purposes of this class, we are interested in the Balmer series of H lines that originate from the n=2 level and occur in the optical part of the spectrum.
The graph in the textbook, Fig. 8.7, shows the relative number of atoms in the n=2 level, as a function of temperature.
The graph shows that very high temperatures are required for significant population of the n=2 level. Why then does the strength of the Balmer lines get weaker as the temperature increases? ie. Why are the Balmer lines so weak in O stars?
The answer is, the H atoms are mostly ionized………
Saha EquationNII = 2 A (kT)3/2 ZII e –h/kT
___ _________________
NI ne ZI
Where=
NI = number of neutral atoms
NII = number of ionized atoms
ZI = partition function for neutral atom (= 2 for H in gd. State)
ZII = partition function for ionized atom (=1 for H)
ne= electron number density ( e/m3)
A is a constant = 2 me/h3
Application of the Saha Equation
The graph in the textbook, Fig. 8.8, shows the fraction of atoms that are ionized as a function of temperature.
The graph shows that H becomes ionized at relatively low temperatures corresponding to ~ 104 K.
The combination of these two equations; Boltzmann and Saha is illustrated in Fig. 8.9 which shows that the population of the n=2 level is sharply peaked at ~ 104 K. So, the Balmer absorption lines are strong only in the atmospheres of stars with temperatures ~ 104 K. (see previous figure on slide 17).
The Harvard Spectral sequenceO, B, A, F, G, K, M
The spectral sequence is a much more precise way to measure the temperature of stars than the B-V color index, discussed previously.
The reason is because the B-V color index is directly affected by the stellar absorption lines which leads to an “incorrect” temperature, since stars are not perfect blackbodies after all.
The Hertzsprung-Russell Diagram
This diagram requires a measurement of the distanceto each star ( to get Mv) and a spectrum for each star( to get the Spectral type ).
Two More Methods to Measure Distances to Stars
Spectroscopic Parallaxes
The name of this method is a bit deceiving as it incorrectlyimplies some measure of stellar parallax. Actually, what happensis we obtain a spectrum for a star of unknown distance. We use thespectrum to determine the spectral type, which locates it on the x-axis of the H-R diagram. Now draw a line up to the main sequence,and continue it horizontally to determine it’s absolute magnitude, Mv. The absolute magnitude combined with the apparent magnitude, mv, allows the distance to be determined using
mv - Mv = 5 log d - 5.
The Second Method is calledMain Sequence Fitting
In this method, we plot a B-V color vs. apparent magnitudediagram for a cluster of stars of unknown distance.
This graph is then shifted vertically over a copy of a calibrated H-R diagram until the two main sequencesoverlap.
The difference between the cluster apparent magnitudeand the calibrated H-R diagram absolute magnitude isthe distance modulus mv - Mv which yields the distance tothe cluster using mv - Mv = 5 log d - 5.
The M-K Luminosity Classification scheme