Phys 1810: Lecture 9

1. Phys 1810: Lecture 9 • Next Class • 21 cm emission lines 18.4 • synchrotron emission • Doppler Shift 3.5, Box 3-3, 4.5, • Telescopes 5.2, 5.3, “seeing” in 5.4, 5.5-5.7 • CCDs, photometry, spectroscopy • Light gathering power, Resolving Power, Diffraction Limit, Seeing

2. Please join us this week, and the first Thursday of every month, rain or shine. Friends and family are welcome too! October 2 at 7:30 pm Meet at Lockhart Planetarium (University College Room 394) Also this month: Oct 1-30:  An exhibit of astronomy images in Degrees Café. October 8:  total lunar eclipse! October 14: Astronomy in the restaurant – The Tallest Poppy at 7:30 PM.  A panel presentation & opportunity for the public to ask questions. Special guest Professor Ken Freeman (Australian National University). Topic: dark matter – The stuff that makes up 90% of the matter in the universe. October 23: partial solar eclipse!

3. Stars: Their Characteristics • Luminosity (L): The total energy radiated per second, at all wavelengths. • L = surface area * flux • Surface area of a sphere is T== Surface Temperature  Luminosity is proportional to the radius squared times surface temperature to the 4th power.

4. Stars: Why Temperature is useful. • Notice that if we know the temperature of a star, then if we know the radius, we can calculate the luminosity. • Alternatively, if we know the temperature and the luminosity we can determine the radius. 

6. The Interaction of light and matter. summary Recall column • Photons (γ == gamma) • Individual packet of EM energy that makes up EM radiation • γ & matter interact creating spectra. • Spectra used to assess • T (blackbody curve type spectrum) • processes that produce light or absorb it (i.e. what is going on) (Animation)

7. Spectra Kirchhoff’s Laws summary Recall column • 3 empirical laws • Hot opaque body -> continuous spectrum • Cooler transparent gas between source& observer -> absorption line spectrum • Diffuse, transparent gas -> emission line spectrum

8. Spectra • Our sun and other stars have an atmosphere. Imagine that you are in a spaceship far above the Earth’s atmosphere. Which of the following spectra would you observe when analyzing sunlight? • Continuum rainbow-like spectrum • Dark line absorption spectrum • Bright line emission spectrum

9. Spectral Finger Prints Solar Spectrum • Note emission lines for lab spectrum of iron are at same λs of absorption lines of iron in  • Can use line spectra to determine chemical elements in object.

10. Interaction of Light and Matter: How are line spectra created? γs of light interact with atoms & molecules. • Atoms consist of: • Electrons (negative charge) == e- • Nuclei (balance charge of e-) • Protons (positive) • Neutrons (neutral) • Molecules: group of 2 or more atoms.

11. Interaction of Light and Matter • Hydrogen == H: simplest atom. • 1 e- & 1 proton. • Classical picture: e- in an orbit. • Contemporary picture: e- as a cloud. • Orbits are really energy levels. • E == energy

12. Interaction of Light and Matter Hydrogen Atom Energy Levels • Every chemical element has its own specific set of E levels. • Each E level is associated with a λ.

13. Interaction of Light and Matter Creating spectral lines at visible wavelengths • specific (quantized) E levels. Level with lowest E is ground state. • How does e- get excited? • By interactions between γs & matter.

14. Interaction of Light and Matter: Creating spectral lines at visible wavelengths • The e- can shift between E levels by absorption & emission of γs.

15. Interaction of Light and Matter Creating spectral lines at visible wavelengths Absorption: • If γ’s E is not matched to any E level then γ passes by atom. The atom is unchanged.

16. Interaction of Light and Matter Creating spectral lines at visible wavelengths Absorption: 2. If γ’s E matches E needed to cause an e- to jump to a larger E level, then atom absorbs γ (i.e. absorbs E) & e- jumps to that E level. The atom is now in an excited state.

17. Interaction of Light and Matter Creating spectral lines at visible wavelengths Absorption: 3. If γ’s E is larger than any jump within atom, then atom absorbs E, the γ disappears, & an e- (or more) are kicked out of the atom creating an ion. (In an ion the charge is not balanced.) The atom is ionized.

18. Interaction of Light and Matter • e-s originally at these E levels in  atmosphere are kicked into excited or ionized states. •  wavelengths of absorption lines are identical to wavelengths of emission lines.

19. Interaction of Light and Matter Creating spectral lines at visible wavelengths Emission: • γs can be emitted spontaneously when an e- falls back down to lower E levels. • An atom can be excited or ionized. An ejected e- can subsequently be recaptured. (animation)

20. Interaction of Light and Matter Creating spectral lines at visible wavelengths Emission: • These e-s cascade through different E levels, generating γs • Bottom path  E level called “H α ” & glows red • (α== “alpha”)

21. Interaction of Light and Matter Creating spectral lines at visible wavelengths Emission: The Orion Nebula David Malin • Clouds of gas that glow due to this process have a few names: • Emission nebulae • H II regions • H α regions • If very bright, then pinker. • Ionizing γs come from hot stars.

22. Spectral Finger Print Hydrogen Atom Energy Levels • Each chemical element has its own “finger print” of lines. • The # of lines for one element depends only on the # of E levels in its atom. • The more elements in a star, the more lines in the star’s spectrum.

23. Spectral Finger Print • The strength of absorption lines gives the # of atoms of that element in the gas. • Comparison of strengths of absorption lines of different elements in the gas gives • Density • Temperature • Can get these characteristics for outer layers of star from its absorption line spectrum.

24. What can we do with spectral information? Study activity on the sun! • the Sun in extreme ultraviolet light (Solar Dynamics Observatory.) • false-color image shows emission from highly ionized iron atoms. • Loops and arcs trace the glowing plasma suspended in magnetic fields above solar active regions.

25. What can we do with this information? • Consider stars...

26. Spectral Finger Print • What can we do with this information? • If 2 stars have the same elements, same density, and same temperature then they have the same intrinsic luminosity. • If they have the same intrinsic luminosity we can use their apparent brightnesses to derive their relative distances using the Inverse Square Brightness Law!

27. Instrumentation for observing spectra: • spectrograph on optical telescope. • use diffraction grating to disperse light. Spectra at other λ: e.g. ν of radio receiver

28. Neutral Atomic Hydrogen emissionH I: 21 cm emission

29. Spin Flip Transition

31. 4. Radio Continuum Emission. Nick Strobel a) Synchrotron radiation: an e- spirals around a magnetic field line.

32. 4. Radio Continuum Emission. Nick Strobel b) Thermal radiation generated when e- accelerates near p+. (Define accel. includes change in direction.)  "free-free" emission, "bremsstrahlung” or “braking”.

33. W3/4/5 Region Our Milky Way Galaxy: ionized hydrogen around hot stars (thermal). Visible Phenomena Optical image courtesy of Charles Dyer

34. W3/4/5 Region Our Milky Way Galaxy in Radio Radiation: The Interstellar Medium ISM 21 cm radio data: Canadian Galactic Plane Survey atomic hydrogen gas (spin flip transition). Imaging the Invisible – Radio Radiation Following images by Jayanne English, Russ Taylor and Tom Landecker using the Dominion Radio Astrophysical Observatory.

35. W3/4/5 Region Our Milky Way Galaxy in Radio Radiation: The Interstellar Medium Include infrared data (thermal) from the IRAS satellite. Imaging the Invisible – Radio and IR Radiation

36. W3/4/5 Region Our Milky Way Galaxy in Radio Continuum Radiation: Imaging the Invisible – Radio and IR Radiation • Heated dust infrared data (thermal) plus e-s moving in B fields • synchrotron in supernova remnants, galaxy cores, ISM. • free-free in ionized shells.

37. Black body radiation (continuum emission) • Spectral line emission and absorption • Spin-flip transition emission • Radio continuum emission Review of Processes Producing Radiation

38. Dopper Shift • How can we use spectral lines? • What properties of objects can we measure?

40. Recall continuum didn’t move – only lines moved. • Continuum when plotted is a b.b. curve. • For peak of the black body curve to change colour, star would need to travel at least 10,000 km/s. Within our Milky Way Galaxy most stars orbit at a speed of 220 km/s. Even nearby galaxies – which I am showing detailed images of - are moving at only a few thousand km/s.

41. What your eye sees for colour. • mathematically defined colour space (CIE) of colour perception. • λ associated for colours perceived by humans (on the outer edge of the shape). Change in colour in 50 nm. Say we observe at 500 nm rest wavelength.

43. Doppler Shift: • If an object is moving towards you, you will observe its spectral lines are shorter in wavelength. • Analyzing a star’s spectral lines will tell us about its density, surface temperature, rotation, chemical composition but NOT its transverse (side-to-side) motion. • We can use the velocity from the Doppler shift of an object in orbit to measure the mass of the object it is orbiting.

44. Stars in our Milky Way Galaxy