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Gaseous Phases in Galaxies

Gaseous Phases in Galaxies. Uli Klein Univ. Bonn. 1. Introduction. 5. Hot gas. 2. Atomic gas. 6. Heating and cooling. 3. Molecular gas. 7. Galactic Winds. 8. Gas mass and  b. 4. Dust. Guess what and where!. Gas phases. ... in the LMC. 1. Introduction. Interstellar cycle:.

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Gaseous Phases in Galaxies

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  1. Gaseous Phases in Galaxies Uli Klein Univ. Bonn 1. Introduction 5. Hot gas 2. Atomic gas 6. Heating and cooling 3. Molecular gas 7. Galactic Winds 8. Gas mass and b 4. Dust Cetraro, Giugno 3 - 7, 2002

  2. Guess what and where! Cetraro, Giugno 3 - 7, 2002

  3. Gas phases ... ... in the LMC Cetraro, Giugno 3 - 7, 2002

  4. 1. Introduction Interstellar cycle: BBNS : 76% H, 24%He, 10-43He, 10-42H, 10-9 - 10-10Li, Be today : 66% H, 32% He, 2% ‘metals‘ Cetraro, Giugno 3 - 7, 2002

  5. Gas phases: Phase n [cm-3] T [K] fV [%] fm [%] h [pc] Tracer molecular 102 ··· 105 10 ··· 50 < 1 ~20 ~70 CO cold neutral 40 ··· 8050 ··· 2002 ··· 4 ~40 ~140 HI (absorpt.) warm neutral 0.1 ··· 0.65500 ··· 8500~30 ~30 ~400 HI (emission) warm ionized ~0.2~8000~20 ~10 ~900 H radio hot ionized 10-3 ··· 10-2 105 ··· 107 ~50 ~1  1000 [OVI] X-rays Cetraro, Giugno 3 - 7, 2002

  6. ISM consists of different gas phases, i.e. components with different temperatures and pressures. Most of them are in mutual pressure equilibrium: P = n  k  T [P] = dyn cm-2 or P/k = n  T [P/k] = K cm-3 Molecular clouds and dust do not participate in pressure equilibrium. Molecular clouds are self-gravitating  behave like stars Dust has fm  1% and T 20 ···30 K Relativistic component: cosmic rays, in pressure equilibrium with the gas, coupled to it via magnetic fields; CRs have fm  10-6 Average energy density of each component participating in pressure equilibrium: ‹u› = 1.6  10-12 erg cm-3 = 1 eV cm-3 Cetraro, Giugno 3 - 7, 2002

  7. Number density of CRs much lower: assume en energy equipartition between particles and magnetic field; then See also lectures by L. Gregorini obtain Lorentz factor from critical frequency (observing frequency), estimate B-field from synchrotron intensity e.g. B = 5 G, vc = 5 GHz  nrel  1.6 ·10-10 cm-3 Cetraro, Giugno 3 - 7, 2002

  8. The meaning of pressure equilibrium: Assumption is justified because the sound crossing time is much larger than the - mean time between SN shocks - recombination time - cooling time More details: see Appendix... Appendix Cetraro, Giugno 3 - 7, 2002

  9. 2. Atomic gas Neutral hydrogen Formation: in the 1st 3 minutes ... (BBNS) 21 cm line radiation of neutral hydrogen at frequency 10 = 1.42040575178(6) GHz hyperfine transition, interaction of electron and nuclear spin magnetic dipole radiation with A10 = 2.86888 ·10-15 s-1 Level population given by spin temperature Tsp: Since h · « k ·Tsp  relative population always 3:1, dominated by collisions Cetraro, Giugno 3 - 7, 2002

  10. Measure Tsp against con-tinuum sources involving on-off and frequency switching technique Towards continuum source: Frequency switching: off position: so that Note: =(), T=Tb() Cetraro, Giugno 3 - 7, 2002

  11. HI mostly optically thin  Total HI mass Kilborn et al. (1999) total dynamical mass: Cetraro, Giugno 3 - 7, 2002

  12. A massive galaxy: Cetraro, Giugno 3 - 7, 2002

  13. A dwarf galaxy (1/100 M31): Not necessarily the true gas distribution ... molecular gas !? Cetraro, Giugno 3 - 7, 2002

  14. LB ~ 0.5  LMW LB ~ 0.06  LMW LB ~ 0.005  LMW Cetraro, Giugno 3 - 7, 2002

  15. Distrubed galaxies: how much mass? IZw 18 HI Bomans et al. (1997) e.g. NGC 4449 - Mtot ~ 2 ·1010 M (?) - MHI ~ 2 ·109 M - heavily disturbed by 109 M companion (DDO 125) - irregular velocity field in centre Hunter et al. (1998) Cetraro, Giugno 3 - 7, 2002

  16. Ionized hydrogen Massive young stars emit photons with   912 Å  ionize surrounding gas; these HII regions emit • recombination lines (Ly, Ly, ... H, H, ... etc.) • free-free radiation radiative transfer (Rayleigh-Jeans approximation for h· « k ·Te): gff = Gaunt factor for free-free emission Te = electron temperature EM = emission measure hence: I  2 for  » 1 I  -0.1 for  « 1 Cetraro, Giugno 3 - 7, 2002

  17. Cetraro, Giugno 3 - 7, 2002

  18. A prime example of free-free absorption: M82 M82 408 MHz Wills et al. (1997) Cetraro, Giugno 3 - 7, 2002

  19. Radio recombination lines: from recombination rates obtain line temperature Tl: Cetraro, Giugno 3 - 7, 2002

  20. Diffuse ionized gas (DIG) DIG found out to large heights above galaxy planes - ‘DIG’ = ‘WIM’: ne ~ 0.02cm-3 T~ 8000 K traced by H - large scale height (Reynolds 1989): where z is measured in pc. Observed H intensity along the l.o.s.: 1 R = 106/4 photons cm-2 s-1 sr-1 = 2.41·10-7 erg cm-2 s-1 by the way....: ne also from - rotation measures RM - pulsar dispersion measures DM Cetraro, Giugno 3 - 7, 2002

  21. Haffner, Reynolds & Tufte Cetraro, Giugno 3 - 7, 2002

  22. Dettmar et al. NGC3109 NGC4700 Cetraro, Giugno 3 - 7, 2002

  23. Problem of ionization: two possible (and necessary) mechanisms! - photo-ionization - shock ionization - CR ionization line ratios [X/H] (X = NII, SII, OI, OIII) indicate mixture of processes  UV photons from HII regions travel large distances out of the plane without being absorbed (and re-emitted) by dust Appendix total energy of > 1042 erg s-1 required exceeds power of all SNe! correspondence with radio continuum & polarization: magnetic fields ‘guide’ ionizing CRs into the halo Association with star-forming activity is obvious  fountains & winds Cetraro, Giugno 3 - 7, 2002

  24. 3. Molecular gas Molecular hydrogen Since about 20 years it is known that hydrogen in the ISM consists at least as much of H2 as of HI!  maps of neutral hydrogen at  = 21 cm yield an incomplete picture! However: direct measurement of H2 difficult; symmetric molecule, lacks permanent dipole moment Ground state 1+: both electrons in lowest orbital Energy spectrum given by - vibration Ev = e (v + ½) - rotation Ev = Bv J(J + 1) - D  J 2(J+1)2 Bv = h/(8·) e vibration frequency D stretching constant Bv rotation constant  moment of inertia Cetraro, Giugno 3 - 7, 2002

  25. Selection rules  radiative and collisional transitions between even (para H2 , J = 0, 2, 4, ...) and odd (ortho H2 , J = 1, 3, 5, ...) levels strictly forbidden. Transitions within each species allowed, in particular electric quadrupole transitions within v = 0, obeying J = 2  = 28 m, T = 512 K  emission only in hot regions (shocks, stellar vicinity) otherwise absorption against (few) bright sources molecular clouds have T  10 ···50 K (cold),  80 ···100 K (warm)  IR emission of H2 not representative for general ISM? Cetraro, Giugno 3 - 7, 2002

  26. Significance of H2: • H2 is the most abundant molecule in the universe • a significant fraction of non-stellar baryonic matter in spiral galaxies is in H2 • H2 is an important coolant of diffuse gas from T  104 K down to T  100 K • H2 cooling influenced structure formation in the early universe • H2 infrared emission traces warm gas, collisionally and/or radiatively excited • H2 promotes all interstellar chemistry H2 in galaxies: pervasive Tk ~ 10 ··· 30 K nH2  1000 cm-3 GMCs Tk ~ 20 K nH2 ~ 10 2 cm-3 dark clouds Tk ~ 10 K nH2 ~ 10 3 ···10 4 cm-3 cores Tk  40 K nH2  10 4 cm-3 Cetraro, Giugno 3 - 7, 2002

  27. Molecular hydrogen is an indispensable ingredient to star formation, hence for the overall fate of the universe (as we witness it now)! Requirement for structure formation early on: cooling time << Hubble time  cooling rate >> expansion rate i.e., mean free path for interaction of particles and photons must be small enough H(t) = dR(t)/dt / R(t) << cool Tegmark et al. (1997) At recombination, i.e. z  1100, MJ  103 ... 106 M(~ globular clusters) H2 controls early structure formation in bottom-upscenario (Tegmark et al. 1997) Cetraro, Giugno 3 - 7, 2002

  28. Formation of molecular hydrogen: Simply by ‘gluing together’ 2 hydrogen atoms? Basically yes, as coll  1010 · (n/cm-3) s  clouds with n  10 ···100 cm-3 this would imply coll  103 yr However, simple 2-atom collisions cannot form H2, since the formation energy (4.5 eV in the ground state) must be expelled. Emission of photon not possible, since the only repulsive state with energy close to zero, the 3+ state, is not radiatively connected to the 1+ state; this would require a change of electronic spin! How, then, dows it work? Cetraro, Giugno 3 - 7, 2002

  29.  Dust as a catalyst! Reaction rate, i.e. rate to hit a dust grain coll = (vH · ng · <d>)-1 plugging in typical values, one arrives at coll  2 ·1012 · (n/cm-3)-1 s vH = velocity of H atoms relative to (much more massive) dust grains ng = number density of hydrogen atoms d = geometric cross section of dust grains  in clouds with n  105 cm-3, a few H atoms will hit a dust grain per year (!); enough to convert all of the H atoms into H2 in a 103 few years!  once there is dust, H2 forms fast (dust has to have Td < 20 K); becomes efficient at nH2 ~ 105 cm-3 Another process: H + H-  H2 + e- about 103 less efficient, however  important in early universe Cetraro, Giugno 3 - 7, 2002

  30. Destruction of molecular hydrogen: Ionization potential of H2 is 15.4 eV (larger than HI)  destruction mostly via photo-dissociation. Selection rules require two-step process for photo-dissociation: (i) upward transition from 1+ ground state to higher bound electronic state, followed by (ii) radiative de-excitation to vibrationally excited state leading to dissociation. H2 simple  process can be calculated; narrow lines related to bound states imply self-shielding of H2 against UV radiation lifetime of H2 in standard interstellar radiation field H2  103 yr, pd = 5·10-11 s-1 unshielded H2  106 yr, pd = 5·10-14 s-1 for columns of 3 ·1020 mol./cm-2 Van Dishoek & Black (1988) Cetraro, Giugno 3 - 7, 2002

  31. Carbon monoxide Molecular hydrogen most important, but most measurements not representative. Second-most abundant molecule: CO, with [CO/H2]  10-4 higher inertia  lower transition frequency  (J = 1  0) = 115.27 GHz ( = 2.6 mm) 5.3 K above ground  (J = 2  1) = 230.54 GHz ( = 1.3 mm) etc. Isotopomeres: 12C16O 13C16O 12C18O 13C18O 12C17O Abundances: 1 1:60 1:240 1:15000 ? Formation and destruction of carbon monoxide: CO mostly from OH + C+ CO + H+ chemically very stable, large ionization potential (14 eV) destruction by photo-dissociation, Ediss. = 11.1 eV  photons with  < 1120 Å, which implies 912 <  < 1120 Å self-shielding much less efficient than in case of H2 ; becomes efficient at NH2 > 1021 mol./cm-2 ; beyond that the main isotope is optically thick Cetraro, Giugno 3 - 7, 2002

  32. Measuring H2 via CO Underlying mechanism: excitation of CO by collisions with H2 For optically thin radiation, e.g. 13CO column density from measured brightness temperature Tb: with 12CO, optically thick, determine Tex: Tb : brightness temperature Tex : excitation temperature Tc : continuum background temperature Measure Tex with 12CO ( » 1) if we know [13CO/CO] in low-density regions  total column density Cetraro, Giugno 3 - 7, 2002

  33. That’s still not NH2 ...! First determinations of NH2 using the virial theorem; stable molecular clouds: v = line width r = radius of cloud So, for a homogeneous cloud: Cetraro, Giugno 3 - 7, 2002

  34. For density distribution (r) ~ r- Measure total CO luminosity of molecular cloud at distance D: or Define Milky Way:XCO = 1.5 ·1020 mol. cm-2 (K km s-1) -1 Once this has been established • measure ICO  NH2 or • measure LCO  Mvir  MH2 (don’t forget to add HI and to correct for helium!) Cetraro, Giugno 3 - 7, 2002

  35. What does this mean? Solomon et al. (1987) • ICO measures (‘counts’) the number of individual clouds within the telescope beam, weighted by their temperatures • Mvir (the total cloud mass) equals the sum of the atomic and molecular gas mass  ICO is a good measure for the H2 column density (or LCO is a good measure for the H2 mass) Guelin & Cernicharo (1987) Tests: measure • LCO, v, r  correlation Mvir  r ·v2? • check extinction vs. measured gas column density: N(HI+2H2) / Av = 1.8 ·1021 cm-2 mag-1 Cetraro, Giugno 3 - 7, 2002

  36. Other methods/checks: Other methods: • FIR & submm emission (Thronson 1986) S ~ NHI + 2 · NH2 • -rays: interaction of CRs with hydrogen nuclei, subsequent 0 decay (Bloemen et al. 1986) I ~ NHI + 2 · NH2 ~ NHI + 2 · XCO · ICO inelastic collision of CR protons with hydrogen, roughly 1/3 of resulting pions are neutral, decaying into two -rays with mean energy of 180 MeV nH  1 cm-3 predicts L   1039 erg s-1, close to what is measured! Cetraro, Giugno 3 - 7, 2002

  37. Other methods: • X-ray absorption: measure NHI and analyse spectrum of soft X-ray emission  2 · NH2 Exercise: decide whether we view NGC253 from ‘above’ or ‘below’....! ... from below! Cetraro, Giugno 3 - 7, 2002

  38. Caveat: XCO depends on • metallicity (C & O abundance, e.g. Wilson 1995) • radiation fields (dissociation) • density (shielding) • angular, hence linear resolution (XCO depends on r and v) • CR heating (Glasgold & Langer 1973) heating by - energetic particles (1 ··· 100 MeV CRs) - hard X-rays ( 0.25 keV) process: H2 + CR  H2+ + e-(~35 eV) + CR primary electrons heat gas by (ionizing or non-ionizing) energy transfer heating rate (Cravens & Dalgarno 1978; van Dishoek & Black 1986): Cetraro, Giugno 3 - 7, 2002

  39. circumstantial evidence: Klein (1999) but: CR flux at E  100 MeV not known in galaxies .... In any case: • high densities, strong excitation, high metallicities : small XCO (e.g. M82, ULIRGS & mergers) • low densities, weak excitation, low metallicities : large XCO (e.g. dwarf galaxies, halo gas) bottom line: detailed case studies indispensable! Cetraro, Giugno 3 - 7, 2002

  40. Examples a normal galaxy ... M51 a dwarf galaxy ... Large Magellanic Cloud! Cetraro, Giugno 3 - 7, 2002

  41. NGC 4214 D = 4.1 Mpc 3 molecular complexes in distinct evolutionary stages • NW : no massive SF yet; excitation process? • Centre : evolved starburst; ISM affected • SE : SF commenced recently; ICO as in NW canonical threshold column density for SF: NHI ~ 1021 cm-2 comparison with HI  above 1021 cm-2 primarily molecular H2 : self-shielding because of high density dissociation of CO in photon-dominated regions (PDRs)  atomic carbon [CI], [CII] [CI] and [CII] are important coolants of the ISM radiative decay of excited states Cetraro, Giugno 3 - 7, 2002

  42. Two contrasting examples: WLM D = 0.9 Mpc: - little SF, weak radiation field & CR flux - XCO  30  XGal (Taylor & Klein 2001) - below 12 + log(O/H) = 7.9 no CO detections of galaxies (Taylor et al. 1998) Cetraro, Giugno 3 - 7, 2002

  43. M 82 D = 3.6 Mpc: - intense SF, strong radiation field and CR flux high gas density, large amount of dust - XCO ~ 0.3  XGal in central region (Weiß 2000) from radiative transfer models; requires many transitions, including isotopomers  true gas distribution - strong spatial variation of XCO - blind use of XCO leads to false results .... Cetraro, Giugno 3 - 7, 2002

  44. Ultra-luminous Infrared Galaxies (ULIRGS): gas densities comparable to stellar mass densities in the centres of elliptical galaxies (Solomon et al. 1995)!! tracers: molecules with high critical densities (HCN, CS, etc.) Cetraro, Giugno 3 - 7, 2002

  45. ‘Measuring’ temperatures and densities Local thermodynamic equilibrium (LTE) and Large Velocity Gradient (LVG) LTE assumes Tkin = Tex = T, i.e. the same temperature everywhere and for all components everything is ‘thermalized’ remember: gu, gl statistical weights column density of optically thin CO then Cetraro, Giugno 3 - 7, 2002

  46. LVG approach: different molecular species may have different excitation temperatures - assumes that optical thinness is provided by turbulence - rotating clouds, spherical symmetry  velocity is a function of distance from centre of a cloud, i.e. V = V0 · r/r0 - this avoids ‘line trapping’, i.e. photons emitted by certain molecular species in certain transition gets absorbed by the same species - assuming the turbulence v » natural line width, then the photons emitted somewhere in the cloud can only interact with nearby molecules, reducing the global problem of photon transport to a local one Cetraro, Giugno 3 - 7, 2002

  47. LVG requires many transitions of a molecule (J = 1  0, 2  1, 3  2, etc.) and its isotopomeres (12C16O 13C16O 12C18O 13C18O) LVG code calculates for given (fixed) input parameters (abundances, velocity gradient, radiation field, beam filling factor) line ratios in the Tkin - nH2 plane least-squares procedure finds the most likely Tkin and nH2 Cetraro, Giugno 3 - 7, 2002

  48. Distribution of molecular gas in M82 Weiß et al. (1999) Cetraro, Giugno 3 - 7, 2002

  49. An effective path length in LVG: L=|dv/dr|-1· v, where v is the observed line width Velocity gradient and CO abundance are input parameters; then Weiß et al. (1999) Cetraro, Giugno 3 - 7, 2002

  50. Direct measurements of H2 Direct observation rendered difficult, owing to lack of dipole moment Measurements with ISO SWS e.g. NGC891 (Valentijn & van der Werf 1999): S(0): J = 2  0 28.2 m S(1): J = 3  1 17.0 m rotational lines, quadrupole transition, 512 K above ground warm component : 150 - 230 K cooler component : 80 - 90 K could amount to 5 - 15 times the HI mass  significant fraction of DM! Cetraro, Giugno 3 - 7, 2002

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