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The Evolution of Molecular Clouds: Formation and Dynamics

Dive into the intricate world of molecular cloud formation and evolution with simulations and scenarios presented by Enrique Vázquez-Semadeni and collaborators from UNAM and abroad. Explore the complex processes that lead to the creation and transformation of these cosmic structures, emphasizing the role of turbulence, gravity, and feedback in shaping molecular clouds. From the initial compression of warm neutral gas to the eventual onset of star formation, this talk delves into the detailed evolutionary pathway of molecular clouds, shedding light on their fundamental properties and behaviors.

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The Evolution of Molecular Clouds: Formation and Dynamics

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  1. Molecular Cloud Formation and Evolution. Scenario and Simulations Enrique Vázquez-Semadeni Centro de Radioastronomía y Astrofísica, UNAM, México

  2. Collaborators: CRyA UNAM: Javier Ballesteros-Paredes Pedro Colín Adriana Gazol Gilberto Gómez Jonathan Heiner ABROAD: Robi Banerjee (Heidelberg) Patrick Hennebelle (ENS, Paris) Katharina Jappsen Jongsoo Kim (KASI, Korea) Ralf Klessen (Heidelberg) Thierry Passot (Obs. Nice) Dongsu Ryu (Chungnam U., Korea)

  3. I. INTRODUCTION

  4. GMCs seem to form from the HI gas

  5. Engargiola et al. 2003: Study of M33 Color image: HI distribution Circles: GMCs GMCs seem to be the “tip of the iceberg” of the gas densitydistribution. They conclude thatGMCs form out of the HI.(See also Blitz+07, PPV; Molinari+13, PPVI)

  6. This talk addresses: • The general evolutionary scenario, • atomic cloud formation:how do the clouds acquire their • mass? • turbulence (at what level)? • GMC formation: • need for self-gravity and global contraction • filamentary structure from gravitational contraction. • destruction: • through stellar feedback. • The numerical tools and database. • Caveat: no time for discussion of magnetic effects.

  7. II. THE EVOLUTIONARY SCENARIO

  8. II.1 Formation of CNM Clouds

  9. WNM n, T, P, -v1 WNM n, T, P, v1 • When a dense cloud forms out of a compression in the WNM, it “automatically” • cools down • acquires mass • acquires turbulence (through TI, NTSI, KHI – Vishniac 1994; Walder & Folini 1998, 2000; Koyama & Inutsuka 2002, 2004; Audit & Hennebelle 2005; Heitsch et al. 2005, 2006; Vázquez-Semadeni et al. 2006). • The compression may be driven by global turbulence, but mostly by large-scale gravitational instabilities: • HI superclouds: <n> ~ 10 cm-3, M ~ 106 -- 4x107 Msun, gravitationally bound (Elmegreen & Elmegreen 1987; Molinari+13).

  10. The system’s structure: ~ cooling length. Adiabatic or phase-transition front. WNM CNM Unstable Vázquez-Semadeni et al. (2006 ApJ, 643, 245).

  11. For stronger compressions and later times, the compressed layer becomes denser, turbulent, and thick,and continuously grows in mass. Adiabatic shocks Phase-transition fronts Thermally unstable, transient gas Ms ~ 2 Audit & Hennebelle 2005

  12. For moderate compressions, the process may be responsible for the formation of thin CNM sheets(Heiles & Troland 2003)at its early stages. Ms = 1.2, tfin ~ 40 Myr Vázquez-Semadeni et al. (2006 ApJ, 643, 245).

  13. The warm-cold transition

  14. log n [cm-3] log T [K] Gas flows from diffuse medium into dense clumps. log P [K cm-3] B [mG] • Dense clump structure (Banerjee+09, MNRAS, 398, 1082). • In a B0 = 1 mG supercritical simulation, with no AD: Cuts through densest point in clump. The boundaries are “phase transition fronts”, not rigid walls.  Clumps accrete.

  15. II.2 Formation of Molecular Clouds

  16. When a CNM cloud forms from a large-scale compression it may involve enough mass to be strongly self-gravitating in this state. • Upon transition to the cold phase, n increases by 100x, T decreases by 100x,  The Jeans mass, MJ ~ n-1/2T3/2, decreases by ~ 104.

  17. Converging inflow setup Lbox Rinf Linflow • Simulations of MC formation and turbulence generation including thermal instability AND self-gravity(Vázquez-Semadeni et al 2007, ApJ 657, 870; see also Heitsch et al. 2008, 2009). • SPH (Gadget) code with sink particles and heating and cooling. • WNM inflow: • n = 1 cm-3 • T = 5000 K • Inflow Mach number in WNM: M = 1.25 (vinf = 9.2 km s-1) • 1% velocity fluctuations.

  18. Run with Lbox = 256 pc; 2.6x107 SPH particles: (Gómez & Vázquez-Semadeni 2013, ApJ, subm., arXiv:1308.6298) WNM in box is initially Jeans-stable. (Mbox ~ 0.01 MJ) Compression cools and compresses the gas. Dense, cold gas soon becomes turbulent and Jeans-unstable. Note: Local fluctuations collapse earlier than whole cloud. Face-on view.

  19. These clouds are never in true virial equilibrium, but appear virialized because of gravitational contraction. SF starts (17.2 Myr) Collapse starts (~11 Myr) |Eg| Apparent virialization: |Eg| ~ 2 Ek Plots for the dense gas (n > 50 cm-3) Ek Eth (Vázquez-Semadeni et al. 2007)

  20. Ekin driven first by inflow, then by gravitational contraction. Lbox = 256 pc, Linf = 112 pc 4.1 Inflow weakens, collapse starts (12.2 Myr) SF starts (17.2 Myr) Turbulence driven by compression, through NTSI, TI and KHI (Walder & Folini1998; Koyama & Inutsuka 2002; Audit & Hennebelle 2005; Heitsch et al. 2005, 2006; Vázquez-Semadeni et al 2006) 2.7 1.4 ~ 0.5 km s-1 (Vázquez-Semadeni et al. 2007)

  21. This scenario implies that • At least the densest, star-forming MCs might be undergoing global gravitational contraction, not be in equilibrium. • Consistent with various recent observational studies (e.g., Hartmann & Burkert 2007; Peretto+2007; Galván-Madrid+09; Schneider+10). Hartmann & Burkert 2007

  22. A more realistic simulation: • No colliding-flow setup. • Avoid the converging-flow “stigma”. • Decaying random turbulence in WNM @ n=3 cm-3. • Clouds form at local shock sites, grow by gravitational accretion. • See J. Heiner’s talk for synthetic HI and CO observations.

  23. Filament formation by gravitational contraction

  24. Because the cloud contains many Jeans masses, its collapse is nearly pressureless. • Collapse proceeds fastest along shortest dimension (Lin+65): Spheroids  Sheets  Filaments • Because collapse of filaments is slower than that of spheres (Toalá+12; Pon+12), spheroidal fluctuations within a filament collapse earlier than rest of the filament.  Start over?  Rest of filament “rains down” onto star-forming clump

  25. Gómez & VS 2013, (arXiv:1308.6298) 2.6 x 107 SPH particles.

  26. Properties of filaments formed in simulation: • Filaments are flow features(not equilibrium objects) where accretion changes from being ~2D to ~1D: cm-2 Gómez & VS 2013, submitted (arXiv:1308.6298). 2 km s-1

  27. Properties of filaments formed by this mechanism: • Accretion flow along filament develops hierarchy of collapse centers and velocity jumps. Gómez & VS 2013, submitted (arXiv:1308.6298).

  28. Properties of filaments formed by this mechanism: Plummer-like radial column density profile: Simulation: Rc ~ 0.3 pc p ~ 2.4 L ~ 15 pc Nc ~ 1022 cm-2 l ~ 40 Msun pc-1 Arzoumanian+11: 0.01 < Rc < 0.08 pc 1.3 < p < 2.4 L ~ 1 -10 pc Nc ~ 1022 cm-2 lrms ~ 30 Msun pc-1 Gómez & VS 2013, submitted (arXiv:1308.6298).

  29. Position-velocity diagram: Asymmetric radial velocity profile at each position. 2-component Gaussian fit: Two-filament appearance due to asymmetry of radial velocity profile. Compare to Hacar+13. Gómez & VS 2013, submitted (arXiv:1308.6298).

  30. II.3 Molecular Cloud Destruction

  31. If clouds are free-falling, one must confront the Zuckerman-Palmer (1974) conundrum: • Free-fall estimate of SFR: • Observed rate is SFRobs ~ 2—3 Msun yr-1; i.e., ~100x lower. • I.e., need to reduce SFR from free-fall value to observed one (~1/100).

  32. Ionization feedback to the rescue • But not by providing support for the clouds. • ... but rather by cloud destruction. • Colín+13 (MNRAS, 435, 1701): simulations of MC formation and evolution using: • AMR code ART + HD (Kravtsov+1997; Kravtsov 2003) • Stellar feedback with • Imposed Salpeter-like stellar IMF. • Crude radiative transfer.

  33. Qualitatively consistent with observations of gas dispersal around clusters • Leisawitz+1989: • Clusters older than ~ 10 Myr do not have more than a few x103 Msun of dense gas within a 25-pc radius. • Surrounding molecular gas receding at ~ 10 km s-1.

  34. - Mayya+2012: - CO, HI and Spitzer study of environment of Westerlund 1: - Region of radius 25 pc contains only a few x103 Msun. Much less than cluster. - Surrounding molecular gas exhibits velocity difference ~ 15 km s-1.

  35. : Density weighted : Volume weighted Virial parameter a evolves: towards unity while driven by gravity. away from unity when feedback starts.  unbinds clouds. Feedback starts

  36. III. THE NUMERICAL DATABASE

  37. 4 types of simulations: • All with: • Cooling • Self-gravity • Resolution: 3.5 orders of magnitude (~250  0.1 pc). 1. Gadget2 SPH simulations: • Lagrangian method • Sink particles • Current manageable resolutions: up to 5x107 particles. 2. ART (AMR) simulations: • Sink particles. • Ionization (and soon SN) feedback. • A realistic stellar IMF.

  38. 3. FLASH simulations • Sink Particles • Include magnetic field and ambipolar diffusion. • Ionization feedback. 4. Fixed grid simulations (Adriana Gazol & Jongsoo Kim): • No sink particles. • Manageable resolution up to 10243. • Uniform resolution useful for medium- and low-density gas studies.

  39. Volume Density PDF Density Spectrum Gazol & Kim 2013, ApJ, 765, 413 Gazol & Kim 2013, ApJ, 723, 482 <M>: mean value of the local Mach number

  40. Column Density PDF Gazol & Kim 2013, ApJ, 765, 413 Lognormal widths as a function of Mrms Isothermal (Buckhart & Lazarian 2012) σ2ln(Σ/Σ0)= AΣln(1 + bΣ2M2) Red : AΣ =0.084, bΣ =12.5 Blue: AΣ =0.081 bΣ =14.29 Mrms: rms Mach number at mean temperature. <M>: mean value of the local Mach number

  41. IV. CONCLUSIONS

  42. The scenario: • GMCs form by convergence of atomic gas. •  GMCs and their clumps accrete in general. • Early stages form: • First thin sheets... • ... then cold HI clump complexes. • Ensemble of CNM clumps is gravitationally unstable and begins contracting. • Contraction needed for both molecule and star formation. • Nearly pressureless contraction naturally proceeds via sheets  filaments  clumps. • SFR increases in time, until sufficient to destroy clouds. • The simulations: • Available for synthetic observations.

  43. THE END

  44. The magnetic case

  45. B B rv rv subcritical dense subcritical diffuse supercritical diffuse supercritical diffuse rv rv subcritical dense supercritical dense If a cloud (i.e., a dense region) is formed by a compression with a component along the field lines,the cloud’smass-to-flux ratioincreases together with its mass(Mestel 1985; Hennebelle & Pérault 2000; Hartmann et al. 2001; Shu et al. 2007; Vázquez-Semadeni et al. 2011). Conclude: the cloud evolves from a subcritical to a supercritical regime. Example: for B=3 mG and n=1 cm-3, a length L > 230 pc is supercritical.

  46. • Magnetic criticality condition (Nakano & Nakamura 1978): • This is very similar to the column density threshold for transition from atomic to molecular gas, N ~ 1021 cm-2 ~ 8 Msun pc-2(Franco & Cox 1986; van Dishoek & Black 1988; van Dishoek & Blake 1998; Hartmann et al. 2001; Bergin et al. 2004; Blitz 2007). • When taking into account the mass of the dense gas only, the clouds evolve so that they are: • subcritical while they are atomic(consistent with observations of atomic gas, e.g., Heiles & Troland 2005) • supercritical when they become molecular(consistent with observations of molecular gas; Bourke et al. 2001; Crutcher+03, +10). • Conclusion: even in ideal MHD, the M/f of a cloud is a continuously evolving quantity.

  47. Converging inflow setup Lbox B Rinf Linflow • Numerical simulations of molecular cloud formation with magnetic fields, self-gravity and sink particles(Vázquez-Semadeni et al. 2011, MNRAS, 414, 2511). • Use FLASH code (AMR, MHD, self-gravity, sink particles). • AD by Duffin & Pudritz (2009). • Similar initial conditions as non-magnetic simulations with GADGET. • Add uniform field in the x-direction. • Optionally add fluctuating field. • Maximum resolution equivalent to 81923. See also Inoue & Inutsuka (2008) for configuration with B perpendicular to compression.

  48. Whole cloud High-N gas Mass-to-flux ratio is highly fluctuating through the cloud, and evolving. Evolution of the mean and 3s values of M/f = m. m = 1.3 m = 0.9 m = 0.7

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