1. Note that the following lectures include animations and PowerPoint effects such as fly-ins and transitions that require you to be in PowerPoint's Slide Show mode (presentation mode).

2. Atoms and Starlight Chapter 7

3. In the last chapter you read how telescopes gather starlight and how spectrographs spread out the light into spectra. Now you are ready to see what all the fuss is about. Here you will find answers to four essential questions: What is an atom? How do atoms interact with light? What kinds of spectra do you see when you look at celestial objects? What can you learn from a star’s spectrum? Guidepost

4. This chapter marks a change in the way you will look at nature. Up to this point, you have been thinking about what you can see with your eyes alone or aided by telescopes. In this chapter, you begin using modern astrophysics to search out secrets that lie beyond what you can see. The analysis of spectra is a powerful technique, and in the chapters that follow you will use spectra to study stars, galaxies, and planets. Guidepost (continued)

5. I. Atoms A. A Model Atom B. Different Kinds of Atoms C. Electron Shells II. The Interaction of Light and Matter A. The Excitation of Atoms B. Radiation from a Heated Object C. Two Radiation Laws Outline

6. III. Stellar Spectra A. The Formation of a Spectrum B. The Balmer Thermometer C. Spectral Classification D. The Composition of the Stars E. The Doppler Effect F. Calculating the Doppler Velocity Outline (continued)

7. The Amazing Power of Starlight Just by analyzing the light received from a star, astronomers can retrieve information about a star’s 1.Total energy output 2.Surface temperature 3.Radius 4.Chemical composition 5.Velocity relative to Earth 6.Rotation period

8. Atomic Structure An atom consists of an atomic nucleus (protons and neutrons) and a cloud of electrons surrounding it. Almost all of the mass is contained in the nucleus, while almost all of the space is occupied by the electron cloud.

9. Nuclear Density If you could fill just a teaspoon with material as dense as the matter in an atomic nucleus, it would weigh ~ 2 billion tons!!

10. Different Kinds of Atoms The kind of atom depends on the number of protons in the nucleus. Helium 4 Different numbers of neutrons ↔ different isotopes Most abundant: Hydrogen (H), with one proton (+ 1 electron) Next: Helium (He), with 2 protons (and 2 neutrons + 2 el.)

11. Electron Orbits Electron orbits in the electron cloud are restricted to very specific radii and energies. r 1, E 1 r 2, E 2 r 3, E 3 These characteristic electron energies are different for each individual element.

12. Atomic Transitions An electron can be kicked into a higher orbit when it absorbs a photon with exactly the right energy. All other photons pass by the atom unabsorbed. E ph = E 4 – E 1 E ph = E 3 – E 1 (Remember that E ph = h*f) Wrong energy The photon is absorbed, and the electron is in an excited state.

13. Color and Temperature Orion Betelgeuse Rigel Stars appear in different colors, from blue (like Rigel) via green / yellow (like our sun) to red (like Betelgeuse). These colors tell us about the star’s temperature.

14. Black Body Radiation (1) The light from a star is usually concentrated in a rather narrow range of wavelengths. The spectrum of a star’s light is approximately a thermal spectrum called a black body spectrum. A perfect black body emitter would not reflect any radiation. Thus the name “black body”.

15. Two Laws of Black Body Radiation 1. The hotter an object is, the more energy it emits: F =  *T 4 where  = Stefan-Boltzmann constant F = Energy Flux = = Energy given off in the form of radiation, per unit time and per unit surface area [J/s/m 2 ]; Energy Flux

16. Two Laws of Black Body Radiation 2. The peak of the black body spectrum shifts towards shorter wavelengths when the temperature increases.  Wien’s displacement law :  max ≈ 3,000,000 nm / T K (where T K is the temperature in Kelvin)

17. Stellar Spectra The spectra of stars are more complicated than pure blackbody spectra. They contain characteristic lines, called absorption lines. With what we have learned about atomic structure, we can now understand how those lines are formed.

18. Kirchhoff’s Laws of Radiation (1) 1.A solid, liquid, or dense gas excited to emit light will radiate at all wavelengths and thus produce a continuous spectrum.

19. Kirchhoff’s Laws of Radiation (2) 2. A low-density gas excited to emit light will do so at specific wavelengths and thus produce an emission spectrum. Light excites electrons in atoms to higher energy states Transition back to lower states emits light at specific frequencies

20. Kirchhoff’s Laws of Radiation (3) 3.If light comprising a continuous spectrum passes through a cool, low-density gas, the result will be an absorption spectrum. Light excites electrons in atoms to higher energy states Frequencies corresponding to the transition energies are absorbed from the continuous spectrum.

21. The Spectra of Stars The inner, dense layers of a star produce a continuous (blackbody) spectrum. Cooler surface layers absorb light at specific frequencies. => Spectra of stars are absorption spectra.

22. Analyzing Absorption Spectra Each element produces a specific set of absorption (and emission) lines. By far the most abundant elements in the Universe Comparing the relative strengths of these sets of lines, we can study the composition of gases.

23. Lines of Hydrogen Most prominent lines in many astronomical objects: Balmer lines of hydrogen

24. The Balmer Lines n = 1 n = 2 n = 4 n = 5 n = 3 HH HH HH The only hydrogen lines in the visible wavelength range Transitions from 2 nd to higher levels of hydrogen 2 nd to 3 rd level = H  (Balmer alpha line) 2 nd to 4 th level = H  (Balmer beta line) …

25. Observations of the H-Alpha Line Emission nebula, dominated by the red H  line

26. Absorption Spectrum Dominated by Balmer Lines Modern spectra are usually recorded digitally and represented as plots of intensity vs. wavelength

27. The Balmer Thermometer Balmer line strength is sensitive to temperature: Almost all hydrogen atoms in the ground state (electrons in the n = 1 orbit) => few transitions from n = 2 => weak Balmer lines Most hydrogen atoms are ionized => weak Balmer lines

28. Measuring the Temperatures of Stars Comparing line strengths, we can measure a star’s surface temperature!

29. Spectral Classification of Stars (1) Temperature Different types of stars show different characteristic sets of absorption lines.

30. Spectral Classification of Stars (2) Mnemonics to remember the spectral sequence: OhOhOhOhOnly BeBeBoy,Bad AAnAnAstronomers FineFForget Girl/GuyGradeGenerally KissKillsKnown MeMeMeMeMnemonics

31. Stellar Spectra O B A F G K M Surface temperature

32. The Composition of Stars From the relative strength of absorption lines (carefully accounting for their temperature dependence), one can infer the composition of stars.

33. The Doppler Effect (1) The light of a moving source is blue/red shifted by  / 0 = v r /c 0 = actual wavelength emitted by the source  Wavelength change due to Doppler effect v r = radial velocity Blue Shift (to higher frequencies) Red Shift (to lower frequencies) vrvr Sound waves always travel at the speed of sound – just like light always travels at the speed of light, independent of the speed of the source of sound or light.

34. The Doppler Effect (2) The Doppler effect allows us to measure the component of the source’s velocity along our line of sight. This component is called radial velocity, v r.

35. The Doppler Effect (2) Example:

36. The Doppler Effect (4) Take  of the H  (Balmer alpha) line:  0 = 656 nm Assume, we observe a star’s spectrum with the H  line at = 658 nm. Then,  = 2 nm. We find   = 0.003 = 3*10 -3 Thus, v r /c = 0.003, or v r = 0.003*300,000 km/s = 900 km/s. The line is red shifted, so the star is receding from us with a radial velocity of 900 km/s.