The Astonishing Slowness of Star Formation

1. The Astonishing Slowness of Star Formation Mordecai-Mark Mac Low Dept. of Astrophysics American Museum of Natural History

2. Collaborators • Miguel A. de Avillez (U. Evora, Portugal) • Javier Ballesteros-Paredes (UNAM Morelia, Mexico) • Dinshaw Balsara (Notre Dame) • Andreas Burkert (MPI für Astronomie, Germany) • Fabian Heitsch (U. Colorado/Boulder) • Jongsoo Kim (Korea Astronomy Observatory) • Ralf Klessen (Astrophys. Inst. Potsdam, Germany) • Volker Ossenkopf (ESTEC, Netherlands) • Michael D. Smith (Armagh Observatory, N. Ireland)

3. The Oddness of Solid Rock • Rock is dense: 2500 kg m-3 • Even water is dense: 1000 kg m-3: • Stars are denser: 105 kg m-3 at center • The average density of the Universe is 10-27 kg m-3 • Even within galaxies, interstellar gas has a density of 10-21 kg m-3, or 1 atom cm-3. • How did galaxies, stars and planets ever form?

4. Gravitational Stability • Criterion for stability of gas against gravitational collapse found by Jeans (1902). • Pressure opposes collapse: sound waves with speed cs must cross region to communicate pressure changes before collapse λJ density ρ

5. Galaxy Formation • Gas and dark matter uniform to 20 ppm when cosmic microwave background emitted. • Denser regions collapse until pressure supports • Further collapse depends on cooling (atomic collisions excite radiation, which escapes). Virgo Consortium

6. What is a star? pressure gravity • Gas collapses under its own gravity • Densities build up until fusion starts in the center. • Resulting thermal pressure opposes gravity fusion • Star is in hydrostatic equilibrium, with pressure balancing gravity

7. Star Formation Rate • What determines the rate of star formation in galaxies? • Free-fall time • Galaxy lifetimes greater than 109 yr. • Yet star formation continues today. • How are starbursts, low surface brightness galaxies different?

8. Observations • Young stars can be identified by surrounding infalling material, models of stellar evolution. • Youngest stars only observed in dense clouds of interstellar gas and dust. • Densities are high enough to shield interior from hard UV radiation from stars, allowing molecules (primarily H2, but also CO, NH3, H2O) to form.

11. Molecular Cloud Lifetimes • Cloud lifetimes estimated by Blitz & Shu (1980) to be around 30 Myr in Milky Way • Locations downstream from spiral arms • Stellar ages associated with clouds • Much shorter lifetimes of 5-10 Myr proposed by Ballesteros-Paredes et al. (1999), Fukui et al. (1998). • stars >10 Myr old not tied with clouds • Cluster ages vs. associated molecular gas • Individual cloud lifetimes vs. ensemble lifetimes

12. Molecular Cloud Kinematics • Molecular spectral line ratioes show cloud temperatures to be of order 10 K, with sound speeds ~0.2 km/s • Line widths are much broader than thermal, corresponding to random motions of order 1-10 km/s, or Mach numbers 5-50. • Strong shocks should be produced, quickly dissipating the kinetic energy.

13. Magnetic Fields • In standard scenario, magnetic fields: • Convert shocks into Alfvén waves (transverse MHD waves), which acts as a lossless spring that stores and returns kinetic energy, allowing observed supersonic motions to persist • Provide magnetohydrostatic support against collapse—ambipolar diffusion (neutral drift through ions) determines time scale for star formation (Mouschovias, Shu, Nakano) • We find magnetic fields either insufficient or unnecessary for these purposes.

14. Decaying Turbulence • Computations with two methods • ZEUS hydro and MHD (Stone & Norman 1992, ApJ Suppl.). Available from the URLzeus.ncsa.uiuc.edu:lca_home_page.html • Smoothed particle hydrodynamics (SPH), using sink particles (Bate et al.) when G0, on a GRAPE 3 special-purpose computer (ancestor of our GRAPE 6 machines) • Periodic, uniform-density, isothermal cube • Gaussian initial velocity perturbations

15. 1283 2563 time time ML 1999

16. 2563 703 193 323 Kinetic Energy Decay SPH hydro ZEUS hydro ML, Klessen, Burkert, Smith (1998, Phys. Rev. Lett.) weak MHD strong MHD

17. Decay Rate • Quantify loss of energy from turbulent, supersonic flow • Measure kinetic energy of boxes driven with constant energy input • Use constant Gaussian driving pattern with narrow range of wavelengths • Vary energy input rate, wavelength, magnetic field strength

18. k = 2 k = 4 k = 8 ML 1999

19. ML 1999 m = mass v = rms velocity k = wavenumber = 2p/ld

20. Jeans length driving length = 2p/k Mach number How fast does turbulence decay? • Mrms >> p in molec. clouds • ld/lJ < 1 needed for support • Turbulence decays in less than a free-fall time in molecular clouds • Observed motions cannot come from initial conditions.

21. Can Turbulence Support Against Gravitational Collapse? • Analytic work (Bonazzola et al., Léorat et al.) suggests that ld < lJ needed for support • Test by adding self-gravity to ZEUS and SPH turbulence models • Zero or decaying turbulence models both collapse efficiently (Klessen & Burkert) • Resolution of cores difficult as collapse continues, so bracket reality with grid, SPH computations

22. Numerical Considerations • Cannot capture core behavior correctly • Bracket with different techniques • Sink particles in SPH: indestructible once formed • Uniform grid: cores can’t collapse, destroyed easily by passing shocks • Magnetic fields diffuse through grid • Minimum number of zones in a Jeans wavelength required to prevent spurious collapse (Heitsch, Mac Low & Klessen 2001)

23. Klessen, Heitsch, ML (2000)

24. Heitsch, ML, Klessen (2000)

25. Rose Center for Earth & Space Images showing star formation and the formation of an HII region from the new Space Show “Are we Alone: The Search For Life”

26. Mbox = 1 Mbox = 1 Klessen, Heitsch, ML (2000)

27. Magnetic fields reduce the fraction of mass in collapsed objects, but do not prevent local collapse. Heitsch, ML, Klessen (2001)

28. Local collapse • Collapse occurs if • When rincreases, smaller regions collapse • Isothermal shocks give • Unless compressed regions are turbulently supported, they collapse locally despite global support

29. Driven vs. Decaying Projected positions of sink particles from SPH models Klessen, Heitsch, Mac Low (2000)

30. Modes of Star Formation • Slow, scattered star formation occurs in regions supported by turbulence due to low densities or high turbulent velocities. Observed in regions like Taurus. • Fast, clustered star formation occurs in regions that are not supported by turbulence, either due to density enhancements or decay of turbulence. Resembles regions like Orion, or starburst knots.

31. What’s driving the turbulence? • Gravitational collapse fails due to fast decay • Protostellar jets and outflows • Most energy deposited outside clouds • Rotational shear of galaxies via magnetic coupling to gas (Sellwood & Balbus 1999) • Probably gives background value (~6 km/s) • Dominant for low surface brightness galaxies, outer regions of normal galaxies?

32. Supernova Driving • In active star-forming galaxies, SN driving dominates other mechanisms • Strength of driving depends on star-formation rate, allowing self-regulation • 3D models (ML, Balsara, Avillez, Kim, 2001, on astro-ph): • Hydro adaptive grid (Avillez 2000) on 3000 x 3000 x 60,000 light year box with galactic disk, clustered, random SNe, and SN rates 1, 6, 10x Galactic value • RIEMANN MHD framework (Balsara 2000)on 600 light year periodic box with SN rate 12x Galactic value

33. Explosions as Bright as Galaxies, Cassiopeia A supernova remnant • (Type II) supernovae occur when massive star fuses all available elements and gravitationally collapses. • Core forms a neutron star or black hole, while outer layers bounce explosively, releasing 1051 ergs of energy 3 light years Chandra X-ray Observatory

34. Simulations of SN-Driving • Avillez (2000) AMR parallel code • vertical stratification, • equilibrium ionization radiative cooling • isolated and clustered SNe (twice galactic rate) • No self-gravity, molecule formation • Show cut through plane of 3D simulation • 0.625 pc resolution in plane (800 x 800 equiv. resol.) • Log of density • 70 Myr • 1 x 1 kpc shown • RIEMANN MHD framework (Balsara 2000)on 200 pc periodic box (1283) with SN rate 12x Galactic value

35. Show disk movie from Avillez

36. Magnetic Pressure Temperature Thermal Pressue Density ML, Balsara, Avillez, Kim

37. Next Steps • Understanding molecular cloud formation • Turbulent compression vs gravitational contraction, including chemistry • Multi-scale computations including both cloud and core formation to capture entire star formation sequence • Modeling the driving of turbulence • Supernova driving vs. shear flows • Are large-scale star formation rates predictable? (empirical answer is yes: Schmidt laws)

38. Overcooling in Galaxy Formation • Far too many dwarf galaxies cool and collapse around galaxies the size of the Milky Way in numerical simulations neglecting star formation, compared to observations (White & Rees 1978, Klypin et al. 1999, Moore et al. 1999). • Ionization by hard UV radiation (λ < 91.2 nm) alone may (Chiu, Gnedin, & Ostriker 2001), or may not (Navarro & Steinmetz 1997) provide enough heating. • Dwarf galaxy wind disruption could be solution (or contribute to it, Scannapieco et al. 2000, 2001ab).

39. Starburst galaxies • When regulation mechanisms fail, star formation rates can be 100x Milky Way’s • The most massive of the newly formed young stars explode as supernovae in only a few Myr. • In starburst galaxies, these supernovae can drive a wind completely out of the galaxy into the surrounding intergalactic gas.

40. Dwarf galaxies form first. Parameter Space • Typical haloes in LCDM with values photo-evaporation (Barkana & Loeb) Lya cooling H2cooling (Ciardi & Ferrara 00, Tegmark 97) Log Mhalo redshift

41. Blowouts From Isolated Dwarfs • 2D, axisymmetric models • Log density shown in color scale • Box sizes of 90,000 x 45,000 light years ML & Ferrara 1999

42. Goals for Study of Cosmological Blowouts • Better understand feedback from early galaxies • Kinetic and thermal energy ejection into intergalactic medium (IGM) • Stop further infall of gas in own halo • Pollution of IGM by ejected heavy elements • Escape of ionizing radiation through bubbles • Calibrate sub-gridscale models for cosmological feedback in large-scale codes

43. t=initial t=collapse t=50 Myr t=90 Myr

44. t=initial t=collapse t=50 Myr t=90 Myr

45. Mechanical Feedback Results • Accreting gas haloes can suppress ejection • With large enough starburst they can themselves be swept away. • Kinetic energy feedback primarily in form of ejected accreting material, not hot gas

46. Star Formation in the Universe • Efficiency and speed of star formation in galaxies determined by the supersonic turbulent motions in the interstellar gas • Turbulence likely driven by combination of supernova explosions and galactic shear • Efficient star formation in young galaxies drives winds that can retard further growth of that galaxy and probably also nearby galaxies.

47. Local Dwarf Models • Mac Low & Ferrara (1999) models: • dwarf disks with constant surface density in hydrostatic equilibrium • Radii from Ferrara & Tolstoy (2000) • potential of DM dominates (softened isotherm. sph.) • Persic, Salucci, Stel (1996) scaling of DM to visible mass • Starburst energy injected at galaxy centers • Excess cooling of hot gas prevented with tracer field • Conduction approximated with mass injection at center • 50 Myr of SN energy input (instantaneous burst) • Low-density IGM for low-z, isolated dwarfs

48. Numerical Methods • ZEUS-3D (Stone & Norman 1992), second-order, Eulerian, artificial viscosity • Ionization equilibrium cooling of ambient gas, with strength dependent on metallicity Z (semi-implicit energy equation) • Density-dependent heating for thermodynamical balance • Tracer field using Yabe & Xiao (1993) transform • Turn off cooling in hot regions to avoid poisoning

49. Isolated Dwarfs • Metals in SN ejecta escape easily • Hot, shocked ejecta have sound speed greater than escape velocity in galaxies up to LMC size • Mass much harder to strip. • In most galaxies, shock “blows out” to IGM before reaching most of ISM • Mass ejected efficiently only from galaxies with baryonic mass < 106 M (“blowaway”)

50. More Realistic Blowouts • Higher pressures and galactic haloes confine blowouts (Silich & Tenorio-Tagle 1998, 2001) • Blowouts in galactic clusters with high-pressure IGM (Murakami & Babul 1999) • Pressure confines blowout • But ram pressure from orbital motion important • Inclusion of Type I SNe (Recchi et al. 2001) • Cosmological blowouts: infall also can limit mass ejection