1. Astrophysics Yr 2 Session 6 Astronomical Spectroscopy

3. Types of spectra

4. The spectral sequence for stars… …a temperature sequence.

5. Z Andromedae spectrum

6. The Orion Nebula & its spectrum

7. Stellar spectra with calibration spectra. Measuring spectral line wavelengths

8. The Doppler effect

9. Energy levels in atoms Principal quantum number n Angular momentum quantum no. l l = 0,1,2,…n-1

10. 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 = 2n 2 electrons

11. Energy of an energy level Convention

12. Energy of an energy level

13. Electron transitions

14. 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.

15. Ionisation terminology A HII region

16. Selection rules for transitions q.m. – conservation of angular momentum l quantum number must change by +/-1 s must not change Rules obeyed → permitted transition Rules broken → forbidden transition

17. Spectral series e.g. Balmer series - hydrogen

18. Sodium (& alkali metals) n = 1, 2; closed shells Spectrum produced by n = 3 electron; Transitions involve n = 4, 5 etc. E +½ - E -½  6Å

19. Sodium term diagram 5889 5895

20. Complex atoms E.g. Helium; 2 electrons → 2 possibilities. 1.One electron stays in n = 1 level. Transitions involve only 2 nd electron & higher levels; → Helium singlet series. 2.Both electrons in higher levels; both take part in transitions; → Helium triplet series.

21. L-S Coupling Electric & spin magnetic fields of electrons interact Greater interaction for higher l values. Spin combinations can enhance, diminish or have no net effect on levels. e.g. two electrons → 3 possibilities – Triplet series.

23. Line profiles A spectral line is produced by a vast population of atoms

24. Saturated lines

25. Line strength – equivalent width

26. Line broadening mechanisms

27. Natural broadening Heisenberg Uncertainty Principle;  E  t  ħ – levels are fuzzy

28. Naturally broadened (Lorentz) profile  t shorter for higher levels →  E larger → line broader

29. Thermal broadening Distribution of radial velocities; Normal or Gaussian v Dop = standard deviation or variance of radial velocity In terms of wavelength (see notes):  Dop = stdv of wavelength distribution.  = - 0

30. 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

31. Take natural log of both sides:

33. Synthetic thermally broadened H  line

34. Pressure broadening Can distinguish between giant & dwarf stars

35. Gas motions; e.g. accretion disc

37. Symbiotic star RX Persei H  line

39. P Cygni

40. The Balmer Jump HI ionisation energy from n = 2 level = 3.4eV. → 3647

41. Molecular spectra

42. A Planetary nebula  1 light year

44. Stellar remnant NGC7207 & its spectrum Forbidden lines due to doubly ionised oxygen

45. Next time: Stellar structure & energy sources