This week at Astro 3303
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- Phebe Garrison
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1 This week at Astro 3303 HW #1 is due today; please pass it in. After it is graded, you can include it in your portfolio HW #2 is now posted; it is due next Tuesday. It includes use of TOPCAT to investigate some datasets. You may work on the homework together but: 1. Be sure to name your collaborators 2. Hand in your own copy of the answers Pick up PE #3
2 Radiation as a messenger I ν,in I ν,out Images One image is worth a 1000 words... Hubble Image One spectrum is worth a 1000 images... Spectra van Kempen et al. (2010)
3 What is the purpose of a telescope? 1. A telescope acts like a light bucket, to gather photons. The bigger a telescope is, the more photons it can catch. Bigger is better => collecting area
4 What is the purpose of a telescope? 2. In addition to gathering light, a telescope allows a more detailed view of the structure of a celestial object and/or to discern the presence of multiple objects. This is called the telescope s ANGULAR RESOLUTION
5 A Telescope s Diffraction Limit The ANGULAR RESOLUTION of a (single) telescope is always limited by its DIFFRACTION LIMIT. Minimum angular separation of two objects that can be seen as separate In radians Remember there are seconds of arc in one radian (a useful number to remember).
6 The seeing of an image The seeing of an image is a measure of its quality or sharpness. Because stars are so far away, they appear as points of light in our images. The seeing then is the angular extent of a star in an image. The seeing is always bigger than either (1) the diffraction limit or (2) the atmospheric seeing, whichever is greater.
7 Electromagnetic spectrum Not all electromagnetic radiation reaches the surface of the Earth; some blocked by atmosphere => Space telescopes (Hubble, Chandra, Spitzer, Compton Gamma Ray Observatory, etc.) Near IR Far 10 microns 1 mm 10cm 10m => 100 microns to 1 mm = Far IR/Submillimeter
8 Telescope characteristics Aperture size (collecting area, diffraction limit) Wavelength/frequency coverage Elevation/transparency of atmosphere Angular resolution/point spread function Field of view Spectral bandwidth Spectral resolution Sampling rate (time domain)
9 Hubble Space Telescope 2.4 meter reflecting telescope Image resolution ~ 0.1 arc second Deployed in low Earth orbit on 25 April 1990 Servicing missions: replace instruments/gyros SM04 completed 2009 Current instruments Advanced Camera for Surveys (ACS: optical - fixed by SM04) Space Telescope Imaging Spectrograph (STIS: UV fixed SM04) Wide Field Camera 3 (WFC3: UV to optical) Cosmic Origins Spectrograph (COS: UV) Near Infrared Camera & Multiobject Spectrograph (NICMOS: infrared)
10 Herschel/Planck to L2 Lagrange point: balance of gravity of Sun-Earth-Moon Herschel
11 JWST: Detecting the first stars 6.5 m telescope optimized for infrared (cooled) Sun-Earth L2 point (orbits Sun) Optical Telescope Element (OTE) Integrated Science Instrument Module (ISIM) Cold, space-facing side Sunshield Spacecraft Bus Warm, Sun-facing side Instruments Near Infrared Camera (NIRCam) Near Infrared Spectrograph (NIRSpec) Mid Infrared Instrument (MIRI) Concept Development Design, Fabrication, Assembly and Test science operations... Formulation Authorization Phase A Phase B Phase C/D Phase E ICR ((PNAR) T-NAR Launch 2018
12 ALMA: Exquisite images of early, dusty galaxies Partners: North America, Europe, East Asia 66 antennas; 54 of 12 meters diameter & 12 of 7 meters diam Operational at millimeter to submillimeter wavelengths Located at 16,500 feet elevation in Chile s Atacama altiplano Partial array science on-going Scheduled for completion in 2014
13 Radio Astronomers need Big Telescopes In radians How big would a radio telescope have to be to have a diffraction limit of 1 arc second at a wavelength of 21 cm? Θ = 1.22 X 21 cm Diameter cm = 1 arcsec/206,265 arcsec/radian Diam (cm) = 1.22 X 21 X 206,265 = 5.2 X 10 6 cm = 52 km (!) How can we possibly build a telescope that big???!
14 Aperture Synthesis or Interferometry Sir Martin Ryle: 1974 Nobel prize in physics Use an array of smaller telescopes to achieve the image detail of a larger one that covers (sparsely) the area of the array.
15 Interferometry Combine information from several widely-spread radio telescopes as if they came from a single dish Resolution will be that of dish whose diameter = largest separation between dishes ( aperture synthesis )
16 Properties of some current and future telescopes See today s handout
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18 Radiation from galaxies In the optical regime, we detect the integrated starlight. I Thermal emission = black body radiation I( )= 2h 3 1 c 2 exp(h /kt) - 1 But galaxies can emit lots of radiation at other wavelengths, both continuum and line emission Typical spectrum of active galaxy, i.e. one with accreting supermassive black hole in its nucleus
19 Bolometric magnitudes Optical astronomers use λ; radio astronomers use. In fact: Apparent brightness In the case of stars, we often use the bolometric magnitude, which is the magnitude integrated over all wavelengths. m bol = m V + B.C. Where the bolometric correction is defined as a function of T eff. B.C ~ 0 for F to G stars => peak in the V-band.
20 The Milky Way viewed edge-on
21 Probably formed in the atmospheres of cool stars Interstellar Dust Mostly observable through infrared emission - very cold < 100 K. Radiates lots of energy - surface area of many small dust particles adds of to very large radiating area Infrared and radio emissions from molecules and dust are efficiently cooling gas in molecular clouds. Whispy nature indicates turbulence in ISM IRAS (infrared) image of infrared cirrus of interstellar dust.
22 Barnard 68 B V Z K H J
23 Photometric observations What we observe is: extincted/fainter absorbed/reddened redshifted projected convolved evolved m obs = m intrinsic + Δm ext + Δm int + Δm K a) Extinction by the Earth s atmosphere (airmass) b) Extinction within the Milky Way (Galactic extinction) Δm ext = Δm b csc b c) Extinction within external galaxy (internal extinction) d) Redshift due to Hubble expansion ( K correction )
24 Interstellar extinction A star of apparent magnitude m will appear to have an apparent magnitude m obs = m + A when viewed through a cloud with A magnitudes of extinction. I = I o e -τ, where τ is the optical depth of the cloud A = m obs - m = 2.5 log 10 (I o /I) = 2.5 log 10 (e -τ ) I o = 2.5 τ log 10 (e) = 2.5 (0.4343)τ = 1.086τ 1 mag of extinction => τ =1 The change in magnitude due to extinction is approx. equal to the optical depth along the l.o.s. Color excess: difference in absorption between two passbands: E(B-V) = (B-V) obs - (B-V) intrinsic = A B A V As a result of galactic extinction, m λ M λ = 5 log d 5 + A λ where d is the distance in pc
25 Corrections to observed magnitudes using the DIRBE instrument on COBE and the all sky altas of IRAS (Schlegel et al
26 A brief overview of stellar evolution Stars of different masses have different lifetimes on the Main Sequence The post-main Sequence evolution of a star depends on its mass. More massive stars have hotter cores; they are able to burn heavier and heavier elements. The most massive stars, during the course of the post-ms evolution, produce the elements up to Fe Post MS evolution is not fully stable, so phases tend to be short, can cause the outer layers to be blown off and are sometimes catastrophic. The phase through which a star finally blows off its outer layers, leaving only its remnant core depends on (1) its mass and (2) how much mass it loses in the post MS phase. The final remnant of a star (white dwarf, neutron star, black hole) depends on (1) its mass and (2) how much mass it loses in the post MS phase.
27 Stars don t shine forever The mass of a star determines how long it will live ( size of the fuel tank and rate of gas consumption ). It also determines how bright it will shine. More massive stars evolve faster ( gas guzzlers ). Mass (Msun) Sp. Type Luminosity (Lsun) MS Lifetime (yrs) 60 O5 800, x B0 50, x B x A x G x K x M x M x High mass stars: where the action is, and on short timescales Low mass stars: lock up baryonic mass for Hubble time or longer
28 Evolution of Low Mass Stars 1. Proto-stellar Collapse (short-lived; unstable) Shines by gravitational energy 2. Main Sequence (long-lived; stable) Shines by controlled, stable hydrogen fusion 3. Post Main Sequence (short-lived; unstable) Succession of phases of fusion a. Red giant: hydrogen burning in shell b. Horizontal branch: He burning in core c. Red Supergiant: He burning in shell 4. Planetary Nebula (short-lived; unstable) Star s outer layers blow off Star therefore loses some of its mass 5. White dwarf (final remnant) Shines due to its heat; eventually fades
29 Evolution of a star like the Sun 2. Stable existence on the Main Sequence Sustained, controlled hydrogen burning the core Energy generation balanced by energy radiated Long-lived phase Along the M.S. L R 2 T 4 but also L M 4 Mass
30 Main Sequence phase 2. Main Sequence (long-lived; stable) Shines by controlled, stable hydrogen fusion photosphere Hydrogen burning in core Star like the Sun: core temp: 15 million K
31 Post-Main Sequence 3. After all the hydrogen in the core is converted into helium, heat continues to leak out, and the core collapses because of gravity. A gas that contracts heats up. A shell around the core is heated until it gets hot enough so that hydrogen burning is ignited in the shell. This new source of heat injects energy into the outer star, which then expands. Because the star s radius increases, its luminosity increases => Moves off the M.S. to higher luminosity
32 photosphere Red Giant Helium burning in core Hydrogen burning shell R ~ 50 R Star like the Sun: core temp: ~100 million K
33 Red Giant Phase 3a. The core contracts until Helium nuclei cannot be packed any closer together => degenerate core Helium ash from the hydrogen burning shell falls onto the core, increasing its mass, density and temperature until the temperature reaches ~100 million degrees K. Helium burning in the core ignites. This happens suddenly: helium flash Helium burning is not a stable process The star becomes a red giant.
34 Evolution of Low Mass Stars 1. Proto-stellar Collapse (short-lived; unstable) Shines by gravitational energy 2. Main Sequence (long-lived; stable) Shines by controlled, stable hydrogen fusion 3. Post Main Sequence (short-lived; unstable) Succession of phases of fusion a. Red giant: hydrogen burning in shell b. Horizontal branch: He burning in core c. Red Supergiant: He burning in shell 4. Planetary Nebula (short-lived; unstable) Star s outer layers blow off Star therefore loses some of its mass 5. White dwarf (final remnant) Shines due to its heat; eventually fades
35 Evolution of High Mass Stars 1. Proto-stellar Collapse (short-lived; unstable) Shines by gravitational energy 2. Main Sequence (long-lived; stable) Shines by controlled, stable hydrogen fusion 3. Post Main Sequence (short-lived; unstable) Succession of phases of fusion a. Red giant: hydrogen burning in shell b. Horizontal branch: He burning in core c. Red Supergiant: He burning in shell 4. Planetary Nebula (short-lived; unstable) Star s outer layers blow off Star therefore loses some of its mass 5. White dwarf (final remnant) Shines due to its heat; eventually fades More and more heavy elements Dispersed over wide area via SN Depending on amount of mass loss, can get supernova White dwarf, neutron star or stellar black hole remnant
36 Evolution of High Mass Stars A high mass star also shines by burning hydrogen when it is on the Main Sequence, but The sequence of fusion reactions (still 4H => He) is different (the C-N-O cycle not the proton-proton chain). Although a high mass star has more hydrogen to begin with, it burns its hydrogen must faster than a low mass star does. High mass stars have shorter Main Sequence lifetimes than low mass stars. The basic sequence of Post Main Sequence Evolution is the same for high and low mass stars, until the star reaches the Red Supergiant Phase. In a low mass star, the carbon-oxygen core cannot burn further. In a high mass star, the core temperature and density are so high that further burning sequences take place. => Produce heavier elements
37 Element production overview
38 Basics of chemical evolution H and He were produce early in the history of the Universe, while all other elements (except for a very small fraction of Li) were produced through nucleosynthesis in stars. Metals are found in very similar (though not identical) proportions in all *s. => the small differences reveal clues about the material from which the *s were made. Z fraction by mass of heavy elements Z ~ 0.02 Most metal poor *s in MW have < 10-5 Z Metal abundance of the ISM gas and of subsequent generations of *s should increase with time (assuming no gas infall from outside) Expect a relation between ages and Z of *s On avg, older *s contain less Fe than younger ones Partially true for Solar neighborhood
39 ZAMS
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42 Color-magnitude diagrams (CMDs) The horizontal axis of the H-R diagram can be represented in different ways: Spectral class Effective surface temperature Color index : difference in brightness measured in separate (wavelength) filters
43 Main Sequence Turnoff: Age Estimator Open clusters: span range of ages
44 CMDs of OCs of different ages
45 Color-magnitude diagrams (CMDs) Also useful as a distance indicator Extragalactic distance scale discussion (later)
46 Main stages of stellar evolution in CMD
47 Critical stages = CMD features Red Giant branch: RGB All stars of low to intermed mass ( M ) Shell H fusion with inactive He core Horizontal branch: HB Phase immediately following RGB Core helium fusion plus shell H fusion Includes RR Lyrae stars Asymptotic giant branch: AGB Can be further divided into E-AGB (early: helium fusion in shell) and TP- AGB (thermally-pulsing: shell fusion episodic). Critical phase for mass-loss Red clump (RC): metal-rich counterpart to HB; good standard candle Post AGB (P-AGB): intermediate stage between AGB and PN for intermediate mass stars (< 2 M ) caused after multiple episodes of dredge-up and so that C is more abundant that O wrt earlier. Blue loop (BL): phase of repeated core He fusion.
48 The method of spectroscopic parallax 1. Observe the star s apparent brightness. 2. Observe the star s spectrum; determine its spectral class and luminosity class. 3. Place the star on the H-R diagram; estimate its luminosity. 4. Use luminosity and apparent brightness to get distance. Also works for star clusters, nearbygalaxies tip of the Red Giant Branch
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52 NGC 5286: CMD
53 NGC 5286: CMD
54 V TO = /- 0.1 (B-V) TO = / NGC 5286: CMD
55 NGC 5286 vs M 3
56 Galaxies: Not a single population Star clusters: assume all stars formed at the same time Galaxies: a complex mixture of stars formed at different times Star formation rate Initial mass function
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