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The Formation of Massive Stars and Star Clusters

Jonathan C. Tan ETH Zurich Christopher F. McKee UC Berkeley Laurent Eyer Geneva Observatory Margaret Kirkland Princeton University. The Formation of Massive Stars and Star Clusters. Outline. 1. Motivation and Open Questions. 2. Overview of physical scales in star and

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The Formation of Massive Stars and Star Clusters

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  1. Jonathan C. Tan ETH Zurich Christopher F. McKee UC Berkeley Laurent Eyer Geneva Observatory Margaret Kirkland Princeton University The Formation of Massive Stars and Star Clusters

  2. Outline 1. Motivation and Open Questions 2. Overview of physical scales in star and star cluster formation 3. The mode of individual star formation in clusters: - Competitive accretion by stars - Formation from Cores 4. Testing models in the Orion Nebula 5. Fast or Slow Star Cluster Formation? The importance of dynamical ejection of massive stars

  3. Notation: Clump -> star cluster GMC Core -> star/binary

  4. Outline 1. Motivation and Open Questions 2. Overview of physical scales in star and star cluster formation 3. The mode of individual star formation in clusters: - Competitive accretion by stars - Formation from Cores 4. Testing models in the Orion Nebula 5. Fast or Slow Star Cluster Formation? The importance of dynamical ejection of massive stars

  5. Motivation Massive stars and star clusters form together. (Gies 1987; de Wit et al. 2004) Most stars are born in relatively rich clusters. (Carpenter 2000; Adams & Myers 2001; Lada & Lada 2003) Total mass in massive stars is small, but they dominate the radiative feedback on the gas. -> Massive stars likely influence the formation of a large fraction of all stars, and maybe planets. -> Massive stars energize the ISM and IGM with their winds, radiation and supernovae; they enrich the gas with heavy elements that set the cooling, thus regulating galactic evolution.

  6. Star clusters appear to be the fundamental unit of star formation (e.g. Carpenter 2000; Lada & Lada 2003)

  7. Some open questions astro-ph/0408491 1. What are the lifetimes and formation mechanisms of GMCs? 2. What, if any, are the triggers for star cluster formation in GMCs? 3. Does star formation in clusters proceed via the collapse of equilibrium cores? 4. What sets the shape and upper limit of the IMF? 5. Are there disks around massive protostars? 6. What sets the properties of binary and planetary companions to massive stars? 7. How are runaway OB stars created? 8. What is the nature of outflows from massive protostars? 9. What feedback processes set the formation efficiency of cores and clumps? 10. How does rotation affect protostellar structure and evolution models?

  8. Why star formation is a difficult problem • wide range of scales (~10 dex in space, time) • complicated physics: (magneto-)hydro, radiation, protostellar evolution, dust and chemical evolution • uncertain initial conditions/boundary conditions • observational constraints are difficult to obtain (nearest massive protostars ~500pc, deeply embedded) There is no computer simulation that has all the physics and resolution necessary to form even a single star

  9. Outline 1. Motivation and Open Questions 2. Overview of physical scales in star and star cluster formation 3. The mode of individual star formation in clusters: - Competitive accretion by stars - Formation from Cores 4. Testing models in the Orion Nebula 5. Fast or Slow Star Cluster Formation? The importance of dynamical ejection of massive stars

  10. R M pressure McKee & Tan (2002) Local examples in the Galaxy:  ~ 0.1 - 1.0 g cm-2 Infrared dark clouds (IRDCs) Star-forming clumps Embedded clusters Pressures in Star-Forming Regions Pressure due to self-gravity: surface density

  11. Surface density vs. cluster mass

  12. =7000 M pc-3 Surface density vs. cluster mass Galactic clumps: Mueller et al. (2002) IRDCs: Carey et al. (2000); Kirkland & Tan, in prep. AV=200 NH=4.2x1023cm-2 =4800 M pc-2 AV=7.5 NH=1.6x1022cm-2 =180 M pc-2 CO clouds AV=1.4 NH=3.0x1021cm-2 =34 M pc-2

  13. Surface density vs. cluster mass Galactic clumps: Mueller et al. (2002) IRDCs: Carey et al. (2000); Kirkland & Tan, in prep. AV=200 Galactic center clusters: e.g. Figer et al. (1999); Kim et al. (2000) Super star clusters: e.g. Gilbert & Graham (2001) De Marchi et al. (1997) Turner et al. (2000)

  14. Outline 1. Motivation and Open Questions 2. Overview of physical scales in star and star cluster formation 3. The mode of individual star formation in clusters: - Competitive accretion by stars - Formation from Cores 4. Testing models in the Orion Nebula 5. Fast or Slow Star Cluster Formation? The importance of dynamical ejection of massive stars

  15. If in equilibrium, then self-gravity is balanced by internal pressure: B-field, turbulence, radiation pressure (thermal P is small) Cores form from this turbulent medium: at any given time there is a small mass fraction in unstable cores. These cores collapse quickly to form individual stars or binaries. Based on SPH simulations with sink particles How do stars form in star clusters? Two different models: Turbulent Fragmentation into CoresCompetitive Accretion McKee & Tan 2003; Bonnell, Vine, & Bate 2004 Vázquez-Semadeni et al. 2004; Schmeja & Klessen 2004 Stars form from “cores”, Mcore~m*, Stars gain most mass by Bondi- that fragment from the clump Hoyle accretion of ambient gas

  16. Objections to Core Model Low and high mass: Turbulence Crowding/interactions Massive stars: Feedback (radiation pressure) Formation timescales Some outflows poorly collimated --> Motivations to consider more dynamic models…

  17. Variable accretion rates Schmeja & Klessen (2004) SPH, periodic boundaries, driven turbulence No B-fields, no feedback, isothermal e.o.s. Sink particle diameter: 560AU . . Mmean~Mjeans/tff Mmax~ (3-50) c3/G

  18. Competitive Accretion & Collisions Competitive Bondi-Hoyle accretion in a star cluster Distribution of gas that eventually joins most massive star Bonnell, Clarke, Bate, Pringle (2001) Bonnell & Bate (2002) Bonnell, Vine, Bate (2004) Stellar collisions Collision timescale Bonnell, Bate, Zinnecker (1998) Radiation pressure problem motivates the radical theory of formation by stellar collisions

  19. Collapsing cores? Disks? Sandell et al. 2003; Chini et al. 2004 Collimated Outflows? Beuther et al. 2002 Chini et al. 2004 Evidence in favor of Turbulent Fragmentation Observational: ~Quiescent cores are seen, both with and without stars (Walsh, Myers, & Burton 2004). Mass function of cores appears similar to stellar IMF (Beuther & Schilke 2004; Motte et al. 2001).

  20. Theoretical: By including B-fields, relatively long-lived cores are seen to form in simulations, even with driven turbulence (Vasquez-Semadeni et al. 2004). SPH sink particle technique, the basis of competitive accretion models, does not adequately resolve Bondi-Hoyle accretion (Krumholz, McKee, Klein 2005). Evidence in favor of Turbulent Fragmentation

  21. Vázquez-Semadeni, Kim, Shadmehri & Ballesteros-Paredes (2004)

  22. log P log  Define effective sound speed Eq. of hydrostatic equilibrium gives log r log r Isothermal case: p=1, c=constant The Conventional Model of Star Formationfrom Near-Equilibrium Cores (Larson 1969; Penston 1969; Shu 1977; Hunter 1977; McLaughlin & Pudritz 1997; McKee & Tan 2003) Self-similar cores -> Polytropic Equation of State Hydrostatic Equilibrium ->

  23. Accretion rates: Collapse from cores at T~10-20 K : log  Works reasonably well for isolated, low- mass star formation (eg Shu, Adams, Lizano 87) log r Form accretion disk: rdisk~ several x  ri ( , typically few %) MHD wind: Accretion rate is constant (p =1, k =2), accelerates (p <1, k <2), or decelerates (p >1, k >2). Can use different boundary conditions, e.g. finite central density and/or finite core mass, bounded by pressure, Ps.

  24. collapse time accretion rate disk size Muench et al. 2002 Protostellar evolution Disk structure Outflows r* m* Properties of cores forming under high pressure Equilibrium cores bounded by P~G2(McKee & Tan 2003) size compact minimum core mass signal speed turbulent

  25. Outline 1. Motivation and Open Questions 2. Overview of physical scales in star and star cluster formation 3. The mode of individual star formation in clusters: - Competitive accretion by stars - Formation from Cores 4. Testing models in the Orion Nebula 5. Fast or Slow Star Cluster Formation? The importance of dynamical ejection of massive stars

  26. Testing the model in Orion 20cm VLA-B Felli et al. (1993) 0.1pc Optical Near IR: VLT-ANTU+ISAAC (ESO)

  27. Becklin & Neugebauer (1967) Kleinmann & Low (1967) 2m point source with no optical counterpart 22m nebula (brightest extra-solar source)

  28. 7 m 8 m 9 m 10 m 11 m 12 m 20 m Image obtained with the MAX camera (MPIA-Heidelberg) on UKIRT by M. Robberto (Robberto, Beckwith, Panagia, et al 2005)

  29. 12.4m IRTF (Gezari, Backman, Werner 1998) BN: 2500-104Lsun =B3-B4=8-12Msun “n” (“L”): not v. luminous or extincted, but has weak radio continuum IRc2: ~1000Lsun : not self-luminous Total L: 8x104Lsun 35 km/s 23 km/s Proper motion BN vs I Plambeck, Wright, Mundy, Looney (1995) Plambeck (p. comm); Rodriguez et al.(2005) To center of Trapezium (0.13pc or 3100yr) Radial vel. +21km/s Scoville et al. (1983)

  30. 12.4m IRTF (Gezari, Backman, Werner 1998) Werner, Dinerstein, Capps (1983) Silicate extinction factor Many of the IR sources are not self-luminous

  31. Chandra-ACIS (Garmire et al 2000) (Getman et al. 2005)

  32. If most luminosity is from a single protostar: L ~ 5.0x104 Lsun protostellar evolution tracks m* ~ 20 Msun m* ~ 3x10-4 Msun/yr . Efficiencies due to outflows ~50% -> Initial core mass ~60 Msun

  33. 12.4m IRTF (Gezari, Backman, Werner 1998) 450m (Wright et al. 1992) 1.4mm (Blake et al. 1996) SiO v=0 J=2-1 (maser) Wright et al. (1996) rcore~0.06 M601/2-1/2pc rdisk~1000 0.02 f*0.71/2 M601/2-1/2AU

  34. Outflows A massive hot protostar may ionize the inner part of the outflow. outflow HII Region disk Density distributions of hydromagnetic outflows (e.g. disk wind, X-wind) approach a common form far from the star or inner disk: collimated wind (Shu et al. 1995; Ostriker 1997; Matzner & McKee 1999). HH30 - HST

  35. Fast blue-shifted OI vr~400 km/s (Taylor, Dyson, Axon, Hughes 1986) BN I

  36. Outflow-Confined HII Regions Fiducial Model: X-wind m* = 20Msun r* = 16Rsun m* = 1x10-4Msun/yr fw = mw/m* = 0.33 fv = vw/vK* = 2.1 Treat sectors independently S = 4 ∫ (2) nw2 r2 dr ri . rc . . 43GHz continuum Menten & Reid, in prep SiO (v=1,2; J=1-0) Greenhill et al. (2003) Tan & McKee (2003) Radio Spectrum: thermal bremsstraulung Radio source “I” appears elongated (0.145”=65AU by <0.085”) at 22GHz (K. Menten, priv. comm)

  37. Fast blue-shifted OI vr~400 km/s (Taylor, Dyson, Axon, Hughes 1986) BN I

  38. Tan 2003 M~60Msun, m*~20Msun, L*~8x104Lsun Testing the model in Orion Near IR (VLT: ANTU+ISAAC)

  39. Outline 1. Motivation and Open Questions 2. Overview of physical scales in star and star cluster formation 3. The mode of individual star formation in clusters: - Competitive accretion by stars - Formation from Cores 4. Testing models in the Orion Nebula 5. Fast or Slow Star Cluster Formation? The importance of dynamical ejection of massive stars

  40. Hypothesis: Star cluster formation is a relatively slow (tform»tff), self-regulated process Motivation: Galactic star-forming clumps often appear round, centrally concentrated and virialized(Shirley et al. 2003) Age spread of young stars in Orion nebula cluster (~3Myr) » tff (~0.2Myr) (Palla & Stahler 1999; Hoogerwerf et al. 2001; Tan 2004) Objections: Turbulence pressure of clump support decays quickly (MacLow et al. 1998; Stone et al. 1998) Feedback from massive stars disperses the gas quickly

  41. Turbulence maintained by protostellar outflows . . Moutflow ~ 0.1 M* voutflow ~ vesc,* 1% coupling of outflow energy Energetics to maintain turbulence: Turbulence maintained with slow star formation rate (tform ~ 40 tff) Norman & Silk 1980 Protostellar outflows are driven by rotating B-fields that are coupled to the protostar and accretion disk(X-winds: Shu et al. 2000; Disk winds: Konigl & Pudritz 2000) Outflow characteristics: bipolar; velocity ~ escape speed; mass loss rate  accretion rate; both low and high mass stars are important Superposition of outflows to form a PROTOCLUSTER WIND (Tan & McKee 2002)

  42. Dispersal of the gas by feedback Galactic clumps: Mueller et al. (2002) IRDCs: Carey et al. (2000); Kirkland & Tan, in prep. AV=200 Galactic center clusters: e.g. Figer et al. (1999); Kim et al. (2000) Super star clusters: e.g. Gilbert & Graham (2001) De Marchi et al. (1997) Turner et al. (2000) Supernova feedback only important after many tff Ionizing feedback: HII region from single O star destroys uniform Galactic clump in ~3x105yr

  43. Dispersal of the gas by feedback Galactic clumps: Mueller et al. (2002) IRDCs: Carey et al. (2000); Kirkland & Tan, in prep. AV=200 Galactic center clusters: e.g. Figer et al. (1999); Kim et al. (2000) Super star clusters: e.g. Gilbert & Graham (2001) De Marchi et al. (1997) Turner et al. (2000) Supernova feedback only important after many tff Ionizing feedback: HII region from single O star destroys uniform Galactic clump in ~3x105yr

  44. Dispersal of the gas by feedback:effect of inhomogeneous, self-gravitating medium Tan & McKee 2000 Model a =1gcm-2 clump as a collection of cores and an intercore medium in virial equilibrium. Form stars near center at a rate such that tform=30tff (50% star formation efficiency in 3Myr) Starburst99 for ionization, stellar winds, radiation pressure Follow evolution of spherically averaged HII region and wind bubble, and motions of cores. Photoevaporation of cores (Bertoldi 1989; Bertoldi & McKee 1990) Mass-loading of wind (Dyson et al. 95)

  45. Test case: single O star with 1049 ionizing photons s-1 HII region Stellar wind bubble

  46. Dispersal of the gas by feedback Galactic clumps: Mueller et al. (2002) IRDCs: Carey et al. (2000); Kirkland & Tan, in prep. AV=200 Galactic center clusters: e.g. Figer et al. (1999); Kim et al. (2000) Super star clusters: e.g. Gilbert & Graham (2001) De Marchi et al. (1997) Turner et al. (2000) Supernova feedback only important after many tff Ionizing feedback: HII region from single O star destroys uniform Galactic clump in ~3x105yr

  47. Feedback destruction times ~Orion ~Arches ~Super star cluster It is relatively difficult to disrupt the clump with stellar feedback This model does not include dynamical ejection of massive stars, which would lengthen the destruction timescale.

  48. Dynamical ejection of massive stars from the Orion Nebula Cluster 4 O and B stars ejected 2.5Myr ago (Hoogerwerf de Bruijne & de Zeeuw 2001) 3 O and B stars in the process of ejection (Tan 2004) ~50% of O and B stars in this cluster are affected by dynamical ejection

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