PHYS 214 Introductory Astrophysics

1. PHYS 214Introductory Astrophysics Instructor: Tracy Webb office: Rutherford Physics Building 217 phone : 514 398 7226 email: webb@physics.mcgill.ca web: www.physics.mcgill.ca/~webb/teaching.shtml office hours: Monday 2-4 PM (or by appointment) Textbook: Astronomy –A Physical Perspective (2nd edition) Marc Kutner (Cambridge) Grade: final exam: 60% mid term: 20% problem sets: 20% Academic Integrity: McGill University values academic integrity; Therefore all students must understand the meaning and consequences of cheating, plagiarism and other academic offences under the Code of Student Conduct and Disciplinary Procedures (see http://www.mcgill.ca/integrity for more information).

2. What is Astrophysics? • study and explanation (quantitative) of everything beyond the earth’s atmosphere (somewhat arbitrary boundary – some would pick a different limit) • therefore a vast enterprise – whole departments in some universities, possible specialization at the BSc level but more often taken up in earnest in graduate school • great symbiosis with many fields of physics – adds to and profits from • classical mechanics – explains planet motion – clean, no friction • atomic and molecular – stellar spectra, interstellar material • nuclear – stellar evolution, burning • particle – cosmic rays (birth of particle physics), dark matter/WIMPs • general relativity - cosmology • astrophysics provides a lab for testing theories and models that would be impractical to test on earth • eg general relativity (gravitational lensing, perihelion of Mercury) • neutrino masses (supernova 1987a gave a limit better than 40 years of lab experiments, solar neutrino experiments established massive neutrinos as real)

3. What is Astrophysics (cont’d) ? • astrophysics relies on results from earth-based experiments • new particles, interaction rates etc have to be incorporated into astrophysical calculations • stellar evolution, nuclear burning requires data from nuclear physics experiments • discovery of new neutrino species has impact on cooling rates in stars and on nucleosynthesis in the Big Bang • picture of the early universe changes when new levels of structure are found in particle physics (eg existence of quarks implies a possible quark-gluon plasma phase of the universe.) • laboratory chemistry provides signatures of elements, molecule production mechanisms etc • recently -- highly dependent on computer models of physical processes - N-Body simulations - Magntetohydrodynamcal simulations - semi-analytic models

4. What will we cover in this course? • essentially everything outside our own solar system • each topic could be a course in itself - so we must be superficial Instrumentation: telescopes, detectors optical, radio, infrared, X-ray, gamma-ray, neutrino astronomy Observing the Universe: atomic theory and spectral lines magnitude scales, stellar distances Physics: black body radiation stellar evolution, interstellar medium galaxies distance scale cosmology

5. Telescopes and other Observing Devices: • ‘experimental’ side of astrophysics is mostly observational (but not passive!) • collect data by examining the object(s) of interest • perform some kind of analysis on the data • very difficult to artificially create a set of controlled conditions like in a classical laboratory experiment (fortunately, a large number of processes are already in progress – just need to find them. Nature has more imagination than most physicists, a bigger budget, and a great deal of time)

6. Optical Telescopes Historical Developments 1610 Galileo - first astronomical use of the refracting telescope (probably a copy of German and Dutch work done earlier) Observed: the Milky Way -- many stars sunspots features on the moon moons of Jupiter phases of Venus strange shape of Saturn from Galileo’s Sidereus Nincius (1610)

7. Telescope Specifications purpose of a telescope is to increase sensitivity and resolution 1. sensitivity depends on light gathering power (LGP) LGP  d2 (d = diameter of objective – lens or mirror) the telescope operates as a ‘light bucket’ with area πd2/4 CFHT 4m telescope compared to the naked eye:

8. LGP is important for studying the distant universe! naked eye can see 6,000 stars 15 cm telescope sees 500,000 30 cm telescope sees galaxies to 200 M light years 5 m 8 B (Universe is ~14 billion years old so 8 billion light years is a fair fraction of the observable Universe.

9. 2. Angular resolution is Resolving Power (RP) ability to distinguish 2 closely spaced objects (eg binary star system, detail in a galaxy) resolution (diffraction limit): qmin = 1.22 l /d radians * q * This is due to light’s wavelike character – it bends around corners (similarly with sound waves and water waves) this is called diffraction Telescope Specifications (cont’d)

10. Diffraction recall basic wave mechanics: apertures can be treated using the techniques of Huygens’ wavelets (infinite number of point sources arranged across the width of an opening, all vibrating in phase at the frequency of the incoming wave) point sources interfere intersection of two crests – constructive interference – bright intersection of crest and trough – destructive interference – dark a screen placed behind a barrier with an aperture will show a diffraction pattern intensity profile the first ‘zero’ is at the angle given by sinq = 1.22 l/d diameter of central spot is ~ l/d (in radians) l = wavelength of light d = diameter of aperture

12. Point Spread Function of the Spitzer Space Telescope (infrared)

13. Consequence for telescopes • plane waves from a point source at great distance (eg a star) • pass through a limiting aperture (lens of refractor or mirror • of a reflector) of diameter d. The spot produced has a minimum • angular diameter of l/d • the star will appear to have a size greater than its intrinsic size • diffraction limits resolution single star binary stars How close can two stars be and still be resolved? no diffraction with diffraction

14. How close can two stars be and still be resolved? answer can depend on statistics - with lots of photons one can perform some kind of fit or on application - need a clear separation if one wants to study only photons from one star of the pair (eg for spectroscopy) qmin ~ l/d radians = 206265 l/d arc seconds typical telescope: d = 1 m l = 500 nm = 500 x 10-9 m qmin = 5 x 10-7rad = 0.1 arc second

15. limits to resolution diffraction is not the whole story; atmospheric turbulence (‘seeing’) limits the angular resolution to 1 or 2 arc seconds large telescopes have the same resolving power as a 10 cm telescope (the large mirrors are needed for light gathering power) this was a prime motivator for sending the Hubble Space Telescope into orbit new developments: discovery that much of the turbulence is local to the telescope better care is now taken to reduce thermal gradients – dome design, electronics isolation etc adaptive optics track a ‘guide star’ and rapidly deform a secondary mirror (eg in a Cassegrain) with computer controlled actuators to keep its image small – actively compensates for the fluctuations

16. HST 2.4m 4m telescope

17. Magnification Power (MP) MP = F/f F = focal length of the objective f = focal length of the eyepiece Magnification changes the size of the image on the eye or detector but does not improve anything. ie image will still be as faint and fuzzy as before

18. different kinds of telescopes: refractors problems with refractors: -chromatic aberration – lenses have a prism effect different wavelengths (colours) are bent different amounts by the glass – images have colour haloes -spherical aberration – lens surfaces are spherical (easier to grind) so the focal point of light depends on the radius at which the ray traversed the lens

19. problems with refractors (cont’d) coma – image of something off-axis is not focused in the same plane as something on-axis (limits field of view laterally so no wide angle viewing) astigmatism -lens is not rotationally symmetric – focal point depends on the azimuthal position at which the light passed through light goes through the lens - two optical-quality surfaces required - highly transparent glass, free from bubbles etc needed - mass proportional to the diameter to the fourth power net result is that large refractors are impractical largest in the world is at Yerkes Observatory (40 inch objective) was built in 1897

20. Reflecting Telescopes • Marin Mersenne proposes the reflector concept • 1663 James Gregory designs a reflecting telescope • 1670 Isaac Newton designs a practical reflector • (and therefore gets the credit) Newtonian Reflector 450plane mirror to avoid the head blocking the light

21. Cassegrain’s Design convex mirror and hole in the main mirror allow an axially symmetric design popular on mid to large size telescopes

22. Advantages of Reflectors • - one optical surface • light does not pass through the objective • can use poor quality glass or other material (metal, epoxy) • no chromatic aberration • support mirror from back and sides • mechanically easier to deal with bigger mirrors • lighter construction – only the surface matters so make the • mirror as thin as stiffness will allow Disadvantages of Reflectors • before invention of aluminized (or silvered) glass • (Foucault – mid 1800s) metal mirrors were used • - subject to temperature fluctuations – deviation from • parabolic shape  bad focus • - subject to tarnishing • secondary mirrors (eg Cassegrain) can cause distortions • coma problems are worse than with refractors

23. Kellner-Schmidt Telescopes reflectors have serious coma problems – objects slightly off-axis are out of focus  field of view is small hybrid design overcomes this patented by Kellner (1910) first constructed by Schmidt (1930) use a spherical mirror instead of a parabola – off-axis light is now focussed but spherical aberration is a problem use a lens to correct the spherical aberration (weak correction is required so lens is thinner and easier to grind/mount than the same diameter lens for a conventional refractor) big advantage: large field of view – use for surveys often called Schmidt cameras largest is the UK Schmidt Telescope (Australia 1973) 1.83 m Edwin Hubble with the 48 inch Schmidt Telescope at Mount Palomar

24. Summary refractor best for detail (potentially the best optics) reflector best for light gathering (study faint objects) Schmidt best for wide angles – statistical studies needing many stars, galaxies

25. Large Reflectors for large reflectors one can observe from the prime focus since the light loss is so small this avoids the need for extra mirrors (more optical headaches) cage support struts cause diffraction patterns on photos light loss due to 1 m diameter cage in a 5 m diameter telescope

26. New Developments in Telescopes Multiple mirrors: old style – monolithic mirror ground and polished from a single glass blank (eg 5 m Hale Telescope on Mt Palomar 20 tonne blank ground to 14.5 tonne mirror correct to 10 nm over the entire face) new style – many thin mirror facets make up the mirror surface mounted with computer-controlled hydraulic actuators to adjust for sag as telescope moves (eg Keck telescopes on Mauna Kea in Hawaii – 10m mirror made from 36 facets of 6 different types Keck 36 mirrors 10m telescope Multiple Mirror Telescope on Mt Hopkins in Arizona six 1.8m mirrors with one common focal point

27. Liquid Mirrors • - old idea recently revived by E. Borra (Laval University) • spin a dish of mercury to get a parabolic reflector • cheap to make • easy to clean • only points at zenith • can be toxic current uses: - satellite tracking - LIDAR - 6 m mirror built for astromony at UBC (prototype for Large Zenith Telescope project)

28. Telescope Arrays build several identical telescopes to gain light collection power of a single large dish (can also be targeted independently if looking at bright objects) eg European Southern Observatory (ESO) Very Large Telescope in Chile four 8 m telescopes to duplicate the area of a single 16 m dish

29. Detectors photographs - emulsions on glass plates (stiff – good for precision) - good for time exposures (see faint objects) - good for permanent records (link up with older observations for long time-base studies) - good for a wide view – lots of information at once but. . . - low quantum efficiency – only ~ 4% of photons striking the plate get recorded (can be improved with image intensifiers) - non-linear – hard to tell from degree of blackening the actual intensity of the light - not digital – difficult to process in a computer

30. Detectors charge coupled devices (CCDs) - solid state devices made from silicon (or gallium arsenide) using integrated circuit technology - array of pixels (picture elements) typically square with up to 6 million elements - highly efficient (~ 100% at red wavelengths – 80% in blue) - photodiode effect: incident photon is converted to electric charge which is stored locally (within the pixel) - can be integrated for extended periods then ‘clocked out’ to an analog-to-digital converter (ADC) – charge gets shifted from pixel to pixel without loss of information (electronic ‘bucket brigade’) - invented at Bell Labs in 1969 - used in digital cameras, video cameras, scanners etc - colour images are obtained using filters main problems are noise and ‘bleeding’

31. bleeding bad pixels

32. Spectrographs • - most knowledge about stars (except position and luminosity) comes from analysing stellar spectra • star radiates like a black body (see later) ie a continuous spectrum of light • atmosphere of star contains various chemicals at low density (low enough • for light to get through) • - atmospheric chemicals absorb specific wavelengths so the continuous spectrum gets holes in it – signatures of the chemicals intensity vs wavelength graph spectrum recorded using a spectrometer

33. spectrometer (spectrograph) prism or grating light from several objects photographic plate or electronic device (eg CCD) collimator to select a specific star (or galaxy)

34. diffraction grating secondary maximum depends on wavelength this spreads out the light into its constituent colours spacing of the slits must be comparable to the wavelength - 10-1000 /mm

35. multi-object spectrograph pass through a grating to get spectra place slits on objects of interest

36. fibre-optic spectrograph 1.. pick out objects in field of view 2. machine a brass plate to fit focal plane and drill holes at the position of the objects 4. get many spectra simultaneously Some spectrographs have robot arms to position the fibres and do not rely on a drilled plate 3. connect optical fibres to focal plane plate and bring other ends to a collimator in front of a spectrograph

37. an example of spectra from my own work: galaxies in a very distant galaxy cluster continuum light emission lines from oxygen and hydrogen

38. most of what we’ve talked about so far applies to optical astronomy but this is only a very small component of observational astrophysics!

39. each wavelength provides measurements of different physics and requires different technology the Milky Way at different wavelengths: radio far-IR near-IR optical X-ray

40. Radio Astronomy • 1931 Karl Jansky at Bell Labs discovers radio emission from outer space while • trying to find cause of static on transatlantic phone calls transmitted by radio • it was discovered that the strongest emission • was from the region of the galactic centre

41. - radio astronomy received a big boost from radar technology developed during WWII • its advantages over optical astronomy are due to the fact that radio waves can traverse • dust, clouds etc both in our atmosphere and in space • can often see what is obscured optically • - can observe 24 h/day rain or shine • - less background problems (especially in third world) • What are we detecting? • - very cold gas • molecular transitions rather than electron level changes in atoms (see later lectures) • emission from energetic electrons • relativistic electrons spiralling in magnetic fields, accelerated in • shock waves, falling into black holes (accretion disks) • What kinds of objects? • growing black holes (Active Galactic Nuclei) • radio jets from massive radio galaxies • supernova remnants • - pulsars (rotating neutron stars)

42. Technology telescopes are large dishes up to ~ 100 m in diameter dishes focus waves onto a sensitive detector element dishes can be made from crude material (eg meshwork to keep weight down) since structure (lack of smoothness) smaller than l/20 is of no consequence eg for 20 cm wavelength, 1 cm holes and bumps are no problem for sub millimetre (newer branch of radio astronomy) one needs 50 mm smoothness the receivers are like audio radios; they can be tuned to pick out a specific wavelength (or frequency – recall that f = c / l )

43. radio astronomy techniques • make intensity maps like in the optical red = radio blue = optical • look for time-related phenomena • - look for periodic signals (eg pulsars) or transients (eg bursts)

44. diffraction in radio astronomy recall that qmin~ l/d lradio = 105lvisible qmin-radio = 105qmin-visible or maybe only 103 to104 since radio dishes are larger by a factor of 10 – 100 one needs dishes 100 km across eg l = 21 cm d = 21 m only gives qmin = 10 mrad = 2062 arc seconds = 0.6o Arecibo dish (Puerto Rico) 305 m diameter operated by NAIC (US National Science Foundation) world’s largest but not-steerable

45. * Lsinq q A B L interferometry • synthesize a large aperture by using more than 1 telescope • - eg 2 telescopes, a distance L= nl apart, look at a single source signals at A and B will be out of phase by the extra time taken to travel the distance Lsinq A B Lsinq

46. as the earth turns, q changes since the source moves across the sky * * * A B A B A B as q changes, so does Lsinq Lsinq = ml A and B are in phase Lsinq = (m+1/2)l A and B are out of phase

47. In general the signals from A and B are somewhere between being in phase and out of phase (ie their relative phase is between 00 and 3600) Adding the signals from A and B gives a signal that oscillates in strength due to the change in interference from constructive to destructive and back change in q between 2 maxima: Lsinq = m l Lsin(q + dq ) = (m + 1) l =Lsinq cosdq + Lcosq sindq = m l + l ~Lsinq + Lcosq dq = m l + l Lcosq dq = l dq = l/ Lcosq now L= n l and for overhead sources cosq ~ 1 dq = l/(n l) = 1/n radians

48. dq = 1/n radians egl = 21 cm L = 21 km dq = 1/n = l /L = 0.21/(21 x 103) = 10–5 radians = 2’’ (like an optical telescope) egl = 30 mm L = 3000 km dq = 10–8 radians = 2 milli-arc-seconds (much better than an optical telescope) • very good angular resolution but only in 1 dimension (like having a CCD • with long, thin pixels). However the earth’s rotation twists the baseline • between telescopes – over 1 day one can synthesize an aperture of diameter L. • (need to trade time for resolution) • many interferometers are arrays . With n dishes one can set up n(n-1)/2 • different baselines simultaneously and gather enormous amounts of information • at once.

49. Interferometer arrays early arrays were linked by cables so necessarily had short baselines in the 1960s atomic clocks and videotape were introduced to allow widely separated dishes to be synchronized off-line (take data independently – radio plus clock signal – then bring tapes to a central facility for analysis) Array examples: VLA (very large array) in New Mexico (27 telescopes – maximum baseline 36 km) MERLIN (multi-element radio linked interferometer) in UK (7 telescopes – maximum baseline 230 km) VLBA (very long baseline array) in US (10 telescopes from Puerto Rico to Hawaii) GMRT (giant metrewave radio telescope) in India (34 dishes – maximum baseline 14 km) Space based interferometer VSOP (Halca and VLBA)

50. Very Large Array Socorro, NM Dominion Astrophysical Radio Observatory Penticton, BC Very Large Baseline Interferometer