astrophysics yr 2 session 6 astronomical spectroscopy
TRANSCRIPT
Astrophysics Yr 2Session 6
Astronomical Spectroscopy
Types of spectra
The spectral sequence for stars…
…a temperature sequence.
Z Andromedae spectrum
The Orion Nebula& its spectrum
Stellar spectra with calibration spectra.
Measuring spectral line wavelengths
The Doppler effect
Energy levels in atomsPrincipal quantum number n
Angular momentum quantumno. l
l = 0,1,2,…n-1
Energy levels in atoms
= 0,+1/-1, +2/-2…+l/-l.
Magnetic quantum number m
Spin quantum number s
= +½ or -½
Spin up Spin down
Pauli Exclusion Principle:n, l, m, s unique to each electron in an atom.
For given n: All l,m,s sublevels full → Closed shell= 2n2 electrons
Energy of an energy level
Convention
Energy of an energy level
Electron transitions
Excitation & ionisation
Collisional/thermal – close encounters with other atoms or free electrons
Photo ionisation/excitation – absorption of photon.For excitation photon energy = energy level difference.For ionisation photon energy> energy of level.
Opposite processes – de-excitation & recombination.
Ionisation terminology
A HII region
Selection rules for transitions
q.m. – conservation of angular momentum
• l quantum number must change by +/-1• s must not change
Rules obeyed → permitted transitionRules broken → forbidden transition
Spectral series
e.g. Balmer series - hydrogen
Sodium (& alkali metals)
n = 1, 2; closed shellsSpectrum produced by n = 3 electron;Transitions involve n = 4, 5 etc.
E+½ - E-½ 6Å
Sodium term diagram
5889
5895
Complex atoms
E.g. Helium; 2 electrons → 2 possibilities.
1. One electron stays in n = 1 level. Transitions involve only 2nd electron & higher levels; → Helium singlet series.
2. Both electrons in higher levels; both take part in transitions; → Helium triplet series.
L-S Coupling
Electric & spin magnetic fields of electrons interact
Greater interaction for higherl values.
Spin combinations can enhance, diminish or have no net effect on levels.
e.g. two electrons → 3 possibilities – Triplet series.
Line profiles
A spectral line is produced by a vast population of atoms
Saturated lines
Line strength – equivalent width
Line broadening mechanisms
Natural broadeningHeisenberg Uncertainty Principle;Et ħ – levels are fuzzy
Naturally broadened (Lorentz) profile
t shorter for higher levels →E larger → line broader
Thermal broadening
Distribution of radial velocities; Normal or Gaussian
vDop = standard deviation or variance of radial velocity
In terms of wavelength (see notes):
Dop = stdv of wavelength distribution.
=-0
Real spectrum; Measure full width of line profile at half peak intensity;
Full Width Half Maximum; FWHM
At line centre (max/min intensity for emission/absorption line)
At ½ max/min intensity
i.e.
f(0)
f(½) FWHM
Take natural log of both sides:
Synthetic thermally broadened H line
Pressure broadening
Can distinguish between giant & dwarf stars
Gas motions; e.g. accretion disc
Symbiotic star RX PerseiH line
P Cygni
The Balmer Jump
HI ionisation energy from n = 2 level = 3.4eV.
→ 3647
Molecular spectra
A Planetary nebula
1 light year
Stellar remnant NGC7207& its spectrum
Forbidden lines due to doubly ionised oxygen
Next time:Stellar structure& energy sources