astrophysics yr 2 session 6 astronomical spectroscopy

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