1 / 42

ISM Lecture 10

ISM Lecture 10. The Three-Phase Interstellar Medium. 10.1 Why necessary?. Two-phase model (see Sect. 8.3) FGH model: Heating by cosmic rays => p.e. effect UV from stars: G ( T ) Cooling by [C II] fine-structure emission: L ( T )

baakir
Download Presentation

ISM Lecture 10

An Image/Link below is provided (as is) to download presentation Download Policy: Content on the Website is provided to you AS IS for your information and personal use and may not be sold / licensed / shared on other websites without getting consent from its author. Content is provided to you AS IS for your information and personal use only. Download presentation by click this link. While downloading, if for some reason you are not able to download a presentation, the publisher may have deleted the file from their server. During download, if you can't get a presentation, the file might be deleted by the publisher.

E N D

Presentation Transcript

  1. ISM Lecture 10 The Three-Phase Interstellar Medium

  2. 10.1 Why necessary? • Two-phase model (see Sect. 8.3) FGH model: • Heating by cosmic rays => p.e. effect UV from stars: G(T) • Cooling by [C II] fine-structure emission: L(T) • Thermal equilibrium: G(T)=L(T) => two stable solutions • Cold Neutral Medium T=50-100 K, nH=10 cm-3 • Warm Neutral Medium T~8000 K, nH=0.2 cm-3 • But….. SN explosions destroy FGH model on timescales <106 yr => need new theory which includes SN activity • Also, vertical scale height of CNM,WNM can only be explained if the medium is turbulent, ~7 km s-1 => something must be stirring up the ISM continuously (mechanical rather than radiative energy) • Note that all theories of ISM are global theories, which try to explain ISM on average, but not locally

  3. Overview of SNR evolution

  4. Why do SN destroy 2-phase ISM? • Frequency S of SN events (not well known) • Probability Q that SN#2 goes off before #1 disappears • Recall • FGH model: no=0.2, vS,fade=10 km/s => tfade=5.2106 yr, Rfade=138 pc => Q=6.(S/10-13) => SN events destroy FGH model in 106 yr => need better theory of ISM!

  5. 10.2 Overview of three-phase model McKee & Ostriker 1977, ApJ Cox 2005, ARAA Tielens Chap. 8.5-8.6 • Q: what is ‘steady-state’ (in global statistical sense) of ISM subject to SN explosions at rate S10-13 pc-3 yr-1? • MO paper is difficult to read; simplified summaries of main ideas given in Appendix

  6. Effect of supernovae on ISM • Overlapping SNRs create “tunnel” system of hot ionized (“coronal”) gas threading through HI clouds, much like the “holes”in a “swiss cheese” • SNR evolves in isolation until it intersects a tunnel and connects to coronal gas; then pressure drops suddenly as SNR “vents” to tunnel system and contributes to pressure of coronal gas • Characterize topology by a “porosity” parameter Q with the fraction of the ISM occupied by coronal gas is f1-exp(-Q), with Q1

  7. Supershells in 21-cm survey of high-latitude gas => large SNR Colomb et al. 1980

  8. Three-phase MO model of ISM • Cold H I clouds • CNM: filling factor f~0.025 • Warm cloud envelopes • WNM f~0.1, ionization by very soft X-rays (?) • WIM f~0.2, ionization by hot B stars • Hot ionized medium • HIM f~0.7 • All three phases are in approximate pressure equilibrium P/k=nT3000 cm-3 K • There is no explicit discussion of dense molecular clouds in MO model, since these have filling factor f10-3

  9. Cross section of a characteristic small cloud • x = ne / n (ionization fraction) • Cold cloud cores give rise to optical absorption lines • Warm neutral medium has low ionization fraction, which is maintained by soft X-ray background (?) • Outer layer is largely ionized by stellar UV background

  10. Small-scale structure of the ISM • A supernova blast wave is expanding into region shown • Small clouds within SNR evaporate

  11. Large-scale structure of the ISM • Only SNRs and large clouds are shown • Altogether 9000 clouds occur in this volume

  12. 10.3 Physical ideas underlying MO modela. Pressure-balance • Various phases are in pressure equilibrium • Pressure is maintained by SNRs • Suppose pressure is initially low  SNRs overlap with internal pressure larger than ISM pressure  pressure in ISM increases until SNRs overlap just when they fade • Pressure balance requires • Note that actual value of Q uncertain especially if magnetic fields taken into account; current estimates range form 0.25-0.75 • This equation can be used to determine PISM

  13. Calculation PISM • Use expressions for Rfade and tfade with vfade=cISM=velocity dispersion ISM in 105 cm/s; no=nH= density into which SNR expands => • PISMnHx1.4mHc2 => • Use

  14. Result PISM • Substitute in above equations => • This result is about a factor of two higher than the average observed insterstellar pressure of 3500 cm-3 K. Above discussion did not include effects of evaporation of clouds which will slow SNR

  15. b. Mass balance • The various phases of the MO model exchange mass. In steady state, these mass flows must balance • Cloud material is lost by “thermal evaporation” due to heat conduction into clouds located within hot young SNR • Cloud material is created in dense cool shell of radiative SNR: clouds make holes in shell and collect material by condensation Cloud evaporation WIM HIM Ionization Radiative SNR WNM Heating Cool dense shell Cold cloud cores Shell fragmentation

  16. Thermal evaporation • What happens to cold cloud inside SNR? Conductive heat flux into cloud Hot medium Cold cloud Mass flow away from cloud If cloud small => evaporates If cloud large => survives + condensation of material on top of it => There is a critical radius for which evaporation = condensation Rcrit strong function of T: RcritnT2 For T>107 K, Rcrit~1020 cm => 30 pc => most clouds evaporate For T<106 K, Rcrit much smaller => less evaporation MO assume cloud size distribution

  17. Supernova density vs. diameter • T higher in young SNRs  • More clouds evaporate in young SNRs  • n(HIM) should be greater in young SNRs than in old (larger) SNRs

  18. c. Energy balance • MO assume that most SN energy leaves system as radiation. This is satisfied if SNRs enter cooling (radiative) phase before they overlap (which they do, both by radiation and cloud evaporation) • A small fraction of SN energy remains as kinetic energy of shells at time of SNR overlap. This source of energy input serves to maintain random motions of clouds. Ultimately, this energy is also converted to radiation following cloud-cloud collisions

  19. d. SummaryInput parameters for MO three-phase model • Supernova rate S • Energy per supernova E51 • Mean density of gas in disk • Ionizing EUV photon production rate per volume • Cloud size distribution and geometries There are uncertainties associated with all these parameters

  20. 10.4 Successes of the MO model • Predicts pressure of the ISM correctly “ab initio” • n vs. age (or R) relation observed for SNRs: younger SNRs are denser • O VI absorption lines: probably produced in cloud evaporation interfaces, where cloud material is heated to T 5105 K

  21. Successes of the MO model (cont’d) • Observed soft X-ray background in rough agreement with T 5105 K expected for cooling SNR gas = HIM • Electron densities in WIM envelopes  0.04 cm–3 agree with observations (pulsar dispersion measures) • Predicted cloud velocity dispersion  7 km s–1 agrees with observations

  22. 10.5 Problems with MO model • MO model does not predict enough WNM; prediction f 0.1 vs. observation f= 0.2 - 0.5 • Cold H I distribution is observed to be much smoother than in MO model • Clouds are more sheet-like than spherical • Type II SNe are not random but occur in clusters of 10 - 40 stars => superbubbles which create galactic fountain (see Sect. 10.6)? • Large fraction of type I SNe occur in disk only and overall SN rate uncertain

  23. Problems with MO Model (cont’d) • For Q 1, expect  3000 supershells covering  50% of disk area on sky – where are they? • O VI absorption lines from cloud interfaces require large energy injection rate – where does it go? • Grain destruction in shocks very effective – how does enough dust survive? • Q 1 requires mean gas density n < 1 cm–3 – is n really that low?

  24. 10.6 Recent developments and modifications of MO model • Scale height of WNM and HIM larger than thought previously • Expect ~20 pc from pressure equilibrium vs. ~200 observed (see following tables) • McKee 1990: This increased “weight” of ISM can be maintained if porosity Q is high, fHIM 0.5 • Filling factors may vary with R and z in the Galaxy, with mostly WIM at large z, not HIM? • External galaxies (e.g., NGC 891) have thick ionized disk

  25. Filling factors with z WIM constrained by new Ha observations by Reynolds (1999) at high latitudes => scale height ~1 kpc. Need O stars whose photons can travel over large distances (kpc) to ionize such a thick layer within MO model

  26. Scale height ISM Blitz 1990

  27. “The Galactic corona” Savage 1995

  28. Recent developments (cont’d) • COBE / FIRAS all-sky survey  surprisingly strong [N II] fine structure lines, consistent with large filling factor and scale height of WIM Fixsen et al. 1999

  29. Three-phase model revisited: disk Purple= mol. clouds Solid green= cold H I clouds Hatched green= warm H I Hatched green on yellow= diffuse warm H II Orange= Hotter gas with OVI Red=Hot gas emitting X-rays Blue= star Top three panels show common (mis)conceptions, bottom panels show features that need to be added. Cox 2005

  30. Three-phase ISM: vertical Same colors as before; Top three panels have difficulties.

  31. Three-phase ISM: global Same colors as before; Top two panels have problems

  32. Summary filling factors in ISM • Little doubt that CNM, WNM, WIM, and HIM are all present in ISM, but filling factors are still controversial • fWNM 0.1 - 0.5 • fWIM 0.2 - 0.5 • fHIM 0.8 - 0.1 MO FGH

  33. 10.6 Gaseous Galactic halo and fountain • Some SNe occur well above Galactic plane and cannot cool efficiently if Thalo 106 K • Can also happen with large OB clusters • This SN energy will be pumped into halo; it can only be removed mechanically  Galactic wind • Thus there may be a Galactic wind with 1 Msun/yr and a halo with Thalo  106 K, nhalo  510–4 cm–3 • Halo would contribute to X-ray background

  34. Observational evidence • High velocity clouds: large number of H I clouds at high latitudes with large velocities  falling into Galaxy? • HVCs may condense from Galactic wind and fall back towards its point of origin • Presence of ions: C3+, Si3+, N4+, … lines seen towards stars in LMC and SMC (see Sect. 6.2)  originate in boundary layers of halo gas and condensing clouds? Or in tenuous gas of T 12000 K? Bregman 1980

  35. 10.7 The local ISM (within  200 pc from the Sun) • Interstellar absorption lines towards nearby stars  very little neutral matter within 100 pc of the Sun, N(H)  1019 H atoms cm–2 • Confirmed by polarization data  n < 0.3 cm–3 • Soft X-ray emission mostly of local origin • O VI absorption lines also indicate that there is a lot of hot gas locally => The Sun is inside a “local bubble” of hot, ionized gas!

  36. Local bubble • Where did bubble come from? SN 107 yr ago? • Is Sun located in unique surroundings? • Is there cold cirrus (as seen by H I, CO, dust) inside bubble? If so, how are they formed and how long do they survive?

  37. Galactic neighborhood 100 pc Orange= Mol. clouds P. Frisch

  38. 10.8 Summary regions in ISM1. H II regions • Traditional H II regions surrounding early type stars • T104 K, n 0.1-104 cm-3, f small • Heated and ionized by photons with hn>13.6 eV, l<912 Å • Cooled by forbidden lines of atoms + ions [O III], [O II], [N II], …. • Observed by optical lines of atoms and ions, radio continuum, radio recombination lines • Coronal gas: very hot, tenuous gas pervading ISM (=HIM) • T3x105 K, n 0.003 cm-3, f 0.6? • Heated by shocks from SN and collisionally ionized • Cooled by adiabatic expansion and X-ray emission • Observed by X-ray emission, optical emission, non-thermal radio emission, UV absorption lines of OVI

  39. H II regions (cont’d) • Warm ionized gas: warm diffuse gas throughout ISM (=WIM) • T8000 K, nH 0.25 cm-3, ne/nH 0.7, f0.2? • Heated and ionized by stray photons from O and B stars l<912 Å • Cooled by forbidden lines of atoms and ions • Observed by broad optical and radio recombination lines (e.g. H166a) and pulsar dispersion measures

  40. 2. H I regions • Cold neutral clouds: H I clouds throughout ISM (=CNM) • T80 K, nH 40 cm-3, f=0.025, ne/nH 10-4 • Heated by UV photons with l>912 Å through p.e. effect • Cooled by [C II] fine-structure emission at 157 mm • Observed by H I 21cm emission + absorption, optical and UV absorption lines of atoms (e.g. Na) • Warm neutral gas: warm diffuse gas throughout ISM (=WNM) • T 8000 K, nH 0.4 cm-3, ne/nH 0.15, f 0.1-0.5? • Heated + partly ionized by soft X-rays + stray photons • Cooled by atomic lines • Observed by H I 21 cm emission

  41. 3. H2 regions • Diffuse molecular clouds: includes translucent and high-latitude clouds • T30-80 K, nH 100-103 cm-3, ne/nH 10-4, f0.01 • Heated by UV photons with l>912 Å through p.e. effect • Cooled by [C II] fine-structure emission • Observed by H I 21 cm emission, CO mm emission, optical + UV absorption lines, IRAS 100 mm • Dense molecular clouds: dark clouds + GMCs • T10-100 K, nH 103-106 cm-3, ne/nH 10-6, f0.0005 • Heated by cosmic rays + newborn stars • Cooled by mm emission from molecules such as CO • Observed by mm lines from molecules, FIR and submm continuum emission from dust Most regions are in approximate pressure equilibrium. Traditional H II regions may be expanding. Dense molecular clouds are self-gravitating

More Related