the origin of stars and their planetary systems
TRANSCRIPT
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The Origin of Stars and Their Planetary Systems
Exploring the Universe with JWST II
24-28 Oct., 2016, Montreal
Ralph E. Pudritz
McMaster University
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Students, postdocs, and collaborators McMaster: Ph.Ds. Mikhail Klassen, Corey Howard, Jason Fiege (U. Manitoba) Alex Cridland and Matt Alessi (planets) Helen Kirk (Banting Fellow, HIA), Bill Harris Samantha Pillsworth (undergrad) Hamburg/Heidelberg: Robi Banerjee, Daniel Seifried (Ph.D.) Ralf Klessen PPVI Andre et al team (Philippe Andre, James de Francesco, Derek
Ward-Thompson, Shu-ichiro Inutsuka, and Jaime Pineda): PPVI Li et al team (Zhi-Yun Li, Jes Jørgensen, Hsien Shang, Ruben
Krasnopolsky, Anaëlle Maury)
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I. Clouds, filaments, and star formation
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Filaments – and the road to cores…
Diffuse Atomic Hydrogen in Milky Way (Canadian Galactic Plane Survey CGPS) near midplane towards Perseus. Extinction map of Orion and Mon
clouds (Cambresy 1998); right - Scuba continuum 850 micron map of 10 pc portion of cloud (Johnstone & Bally 2006)
Shocks and filaments on many scales, in diffuse and self-gravitating gas
Turbulent structure in ISM
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Star forming gas in GMCs: column density structure in molecular clouds (use near IR dust extinction maps)
No star formation: Star formation: Lupus V and Coalsack clouds Taurus and Lupus 1 clouds Lognormal column density Lognormal + power-law Turbulence produces lognormals – gravity the power-law..
Kainulainen +2010
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Column density thresholds for star formation
Left – column density distribution for cores in Aquila (Könyves+2010)
Right – column density distribution for cloud (André+2011, Schneider+2013) Threshold: vertical line Av ~ 8, or N ~ 7 x 1021 cm-2
volume density n ~ 104 cm-3
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Star forming cores in filaments: gravitational instability?
Filaments: lengths 0.1 – 100 pc; masses 1 – 105 Msun , line masses < 1000 Msun pc-1 (Bally+ 1987, Hacar+2013, Kainulainen+2013,2016, .. Review: Andre
+2014) Herschel observations: clouds filamentary down to sub pc (Andre+ 2010, 2014 (PPVI), Menshchikov+2011, Henning+2010..) - Cores strongly associated with filaments > 70%; Polychroni+ 2013)
- GI? White = mass / length (T=10K): m >mcrit = 2 cs
2 / G ~ 16 Msun pc-1
Aquila star forming cloud: Andre+ 2010
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Core mass functions n Well fit by lognormal
distribution.
n Closely resembles the IMF
n Peak at 0.6 Msun
n Compare to peak of IMF:
conversion efficiency of core to stellar mass: ~ 0.2- 0.4 in Aquila Feedback effects?
Core mass function for Aquila cloud (500 starless cores: 0.01-0.1 pc in size) (Könyves+2010, Andre+2010, Andre+ review 2014)
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Filaments and cores live in highly dynamic environments: Velocity structure in filaments
Accretion and filamentary flow: sketch of velocity field around Serpens-South cluster (Kirk et al 2013)
Total velocity dispersion vs central column density of filaments (Arzoumanian et al 2013)
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B fields - channeled accretion flows or infall? (review Hua-bai Li + 2014) Infrared (H band) polarization overlaid on column density map + young stellar cluster
- Orientation – often perp to dense filament? - Field strength ~ 100 µG
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Radiative feedback from massive stars in young clusters – affect filaments? IMF?
Filaments and OB clusters, and HII regions in G10.6-0.4 ( IRAM 30m MAMBO-2 bolometer array + SMA ) - 200 M¤ OB cluster - Ultracompact HII regions: A-E
H.B. Liu et al 2011
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Theory: turbulence, gravity and fragmentation
Turbulence, filaments, and
turbulent fragmentation -Theory; eg. Larson 1981; Elmegreen & Scalo
(2003) -Reviews: eg. MacLow & Klessen 2004; McKee
& Ostriker 2007; Bonnell et al 2007 - Simulations; Porter et al 1994; Vazquez-
Semadeni et al 1995, Bate et al 1995, Klessen & Burkert 2001; Ostriker et al 1999, Padoan et al 2001; Tilley & Pudritz, 2004,2007; Krumholz et al 2007, Federrath et al 2010,…
Shocks dissipate turbulent support as t-1 (eg. Ostriker 2001)
Bonnell et al (2003)
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1. Turbulent fragmentation: - Isothermal gas and supersonic turbulence, 3D simulations: PDF of gas density = log- normal (Padoan & Nordund 1997). Generalize to Include chemistry (Glover+2010). - Star Formation Rate- Integral of density fluctuations above threshold: each collapsing in 1 local free-fall time. Efficiency ε~ 0.3-0.7 - Different models for threshold (review Padoan+ PPVI, (Hennebelle & Chabrier 2011, Fedderath & Klessen 2012, Hopkins )
P(s) = 12πσ s
2exp −(s− so )
2
2σ s2
⎡
⎣⎢
⎤
⎦⎥
s ≡ ln(ρ / ρo )so = −σ s
2 / 2σ s2 = ln[1+ bMs
2 ]Ms =σ v / cs
SFR = εφt
t ff ,ot ff (s)scrit
∞
∫ sP(s)ds
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2. Gravitational instability in filaments 1. Radial force balance for infinite cylinders, constant line mass m ~ ρ r2 ~ const i) Gravitational force: Fg = 2Gm / r ~ r -1
ii) Radial (thermal) pressure gradient for polytropic gas (index γ ) Fp = ρ-1 dP/dr ~ r 1- 2γ -> radial pressure / gravitational force ~ r 2 ( 1- γ)
POINT: γ > 1 stable (pressure wins for small r)
Υ<1 unstable (pressure loses to gravity ultimately) γ=1 constant force ratio -> critical condition. mcrit = 2cs
2 / G = 16 (T/10 K) M¤ pc-1 Values: 2 states, low and high density; γ ~ 0.73, 1.03 (Larson 1985, 2005; Koyama & Inutsuka 2000; low density – Benson & Myers 1989, Evans 1999)…
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Radial density profiles – probing forces on filaments Not observed: isothermal filaments (Ostriker 1964; r -4 ). Herschel: r-1.5 – r-2.5 (Arzoumanian+2011) MHD: Poloidal field supports cylinder, toroidal component *squeezes it*. Fiege & Pudritz 2000a (FP2000ab) Observations
(Matthews & Wilson 2001) r -1.5 – r -2 Accretion: dynamic accretion not static equilibria: r-1.3 - r – 2
(R.Smith+2014; Kirk, Klassen, Pudritz, & Pillsworth 2015) Turbulence: Changes the critical mass to mcrit = 2σ2 / G ; σ2 = cs
2 (1 + M2/3) (FP2000a)
JWST: Column density extinction maps of cores/dense filaments much better than ground based ( Av=50 mags, at high res 10-15”) -> Av = 120 mag -> KAB = 29 mag or 9nJ @ 2µm
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Fragmentation of the “Integral Filament” – ALMA Observations of Orion A (Kainulainen+2016)
Multi-scale fragmentation - Linear theory; periodic - fastest growing mode ~ several times filament width (eg. Larson.., Inutsuka * Miyama,.. Peak at 55,000 AU but growing power at spacings < 16,000 AU - Periodic fragmentation for 55,000 AU? (Fiege & Pudritz 2001b, Fischera & Martin 2012). Jeans type in spherical regions < 16,000 AU?
Surface density of dense cores – contours at 60,80…140 x 1021 cm-2
2 point correlation function of dense cores
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II. Cluster formation and radiative feedback
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IMF in star forming regions
IMFs from 8 star forming regions. Data: Scholz+2012, Alves de Oliveira+2013, Oliveira+ 2009, Sung & Bessell 2010, Luhman+2007, Pena Ramirex+2012, Bayo+2011, Lodieu2013. Compiled by Offner+2014 (ppvi)
Remarkable agreement between different regions: with IMF Chabrier (2005): lognormal - with peak 0.2-0.3 Msun and σ = 0.5 JWST: push down to 1 Mjup by observing in 3-10 µm: - Broad band NIR and MIR
photometry - Classification spectroscopy
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Cluster Formation Including Radiative Feedback Radiative feedback from
massive stars: raises Jeans Mass - filaments don’t fragment - gas drains into primary and its disk (eg. Krumholz et al 2007) - prevent fragmentation out to
1000 AU scales Suppression of brown dwarfs: by factor 4 (Bate 2009): get robust low mass part of IMF…?
2/3TM J ∝
Bate (2009)
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Mass limitation during cluster formation: fragmentation-induced starvation? (Peters, Klessen, Mac Low, & Banerjee 2010)
Collapse into a disk of 1000 Msun clump. Ionization feedback. - Fragments compete for gas in an unstable
disk: muliplicity for O stars… - Radiative heating -> larger Jeans mass - (100 AU resolution)
Disk mass levels out..
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Building Clusters - ionization does not readily erode dense cold filaments (Dale & Bonnell, 2011)….. how does accretion stop? Ionize diffuse gas in the voids instead – filamentary structure has major implications for efficiency of feedback, and therefore SFR Break the filaments with outflows? (eg. Wang et al 2010)
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Radiative feedback and cluster formation in simulated GMCs (Howard, Pudritz, & Harris, 2016, MNRAS)
n Radiative feedback from evolving IMF in cluster sink particles – using ray tracing from sources (Peters + 2010), not FL. Cluster sink particles.
n Clouds idealized: initial n=100 cm-3, box 32- 80 pc, flattop + r-3/2 density profile; low rotation (2 %), Burgers turbulence, Mcl = 106 M¤
n Ionization, thermal heating, radiative momentum..
SEQUENCE in virial parameter α. Observed clouds: α from 0.5 – 5 (Rosolowsky2007, Hernandez & Tan 2015))
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M = 106 Msol, α = 0.25, flatop,..
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Cluster formation: initially poorly bound (α = 3), 106 MSun , GMC
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α and cluster formation
Top: Stellar mass evolution of most massive star clusters – jumps are mergers.. Bottom: clump and star formation efficiencies Unbound model shows efficiencies that agree with obs. Radiative feedback effects negligible…! POINT: low SFE requires initially unbound clouds…
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UV escape fraction from Giant Molecular Clouds (Howard, Pudritz, & Klessen 2016, archive next week)
- Hammer projections of UV ionizing photon fluxes from 106 Msun , α = 3 cloud; - Circles – 10 most luminous clusters, positions projected.
UV escape fraction- highly variable in time with peaks fesc,max ~ 30-35%, average fesc ~ 15%
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III. Massive Star Formation
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Disks around massive stars Keplerian disks observed around 8-30 Msun stars (review: Bertran & De Witt 2013) Disk radii ~ few 1000 AU, masses from 4 to a few 10s of MSun Core-accretion picture favoured? (Myers & Fuller 1997, McLaughlin & Pudritz 1997, McKee & Tan 2002)..
ALMA maps: Keplerian disk around forming O7 star in AFGL 4176 mm1: 25 Msun star, with 12 Msun disk. Top: outflow perp to disk Bottom: CH3CN (J=13-12) (Johnston + 2015)
Best fit is Keplerian disk; velocity peaks of different tracers: (Sanchez-Monge + 2013)
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Massive star formation – how do you make a 150MSun ? Geometry: for spherical infall, Eddington limit for most massive star = 40Msun (Larson & Starrfield 1971, Kahn 1974, Yorke & Krügel 1977, Wolfire & Cassinelli 1987,..)
-1D sims give Mmax ~ 40 Msun independent of core mass > 60 Msun (Kuiper + 2010 ).
Disk accretion -> asymmetry to radiation field (Nakano 1989) - “flashlight” effect (Yorke & Bodenheimer 1997)
- 2D disk sim, multifrequency FLD but unresolved dust sublimation front; - similar mass limit (Yorke & Sonnalter 2002):
Disk in 120 solar mass case quickly destroyed
The problem: need to resolve dust sublimation region!
Yorke & Sonnalter 2002
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Pure FLD vs Hybrid Radiative Feedback (Krumholz+2009 – Kuiper + 2010)
Pure FLD in a 3D AMR code: initial 100 Msun core, r-1.5 density profile, no turbulence.. - R-T instabiltiies @ 41.7kyr, multiple fragmentation
Hyrbid code (ray trace + FLD) - 2D: 120 Msun shown over long time 460kyr – no R-T instabilities
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Radiative Feedback and Massive Star Formation (Klassen, Pudritz, Kuiper, Peters, & Banerjee 2016, ApJ ) - using Hybrid AMR FLASH radiative transfer (Klassen+ 2014)
4 stages: - Gravitational collapse - Disk formation - smooth - Disk instability and massive episodic accretion - Radiative feedback – radiatively driven bubbles
No R-T instability seen (remove ray trace and RT develops, Kuiper+2012) No disk fragmentation up to 50kyr. 100 Msun mode: left face on, right edge on
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Time evolution of central “sink” (massive star) for 3 sims - 30, 100, and 200 Msun cores
1. Initial phase of accretion: - core collapse dM/dt = mo cs
3 / G mo ~ NJ (number of Jeans masses) (Shu,1977 Giichidis + 2012)
2. Disk accretion - epsodic driven by onset of GI. (Vorobyov & Basu 2007). Disk mass levels out…
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Edge- on view of collapse of 200 Msun core zoom in to central region
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Fragmentation and cooling… mostly stable
-1000
0
100015 kyr
-1000
0
100025 kyr
-1000
0
100035 kyr
-1000 0 1000
-1000
0
100040 kyr
-1000 0 1000 -1000 0 1000
0.0 0.6 1.2 1.8 2.4 3.0 3.6 4.2 4.8 5.4 6.0β Stability Parameter
Fragmentation of disks requires both Q< 1, AND rapid cooling (Johnson & Gammie 2003) tcoolΩ = β < 1 Middle: β map ( cooling from flux radiative flux loss): Right column: Q<1 in dark blue: maybe one fragments ~ 1000 AU
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JWST observations..
- No disk fragmentation observed on ~1000 AU scale for t < 50 kyr. Possibly later… - No R-T fallback (But see Rosen+2016) - Primary role of disk instability; drives high accretion rates that are highly episodic.
JWST: fragments in disks? Diffraction limited imaging at NIR? (3.8 and 4.8 microns. Av ~ 150 mag reduced to 4 and 6 mags. Diffraction limited imaging gives 0.15” res… equivalent to 7 fully sampled elements acros 1” = 1000AU @ 1kpc.
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IV. Low mass star formation: disks, outflows, planet formation…
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1. First hydrostatic stellar cores (FHSC) and Class 0 Protostars
- FHSC earliest object forms inner region becomes opaque to radiation (ρc > 10-13 gm cm-3) – collapse becomes adiabatic (Larson 1969). Collapse resumes when when T> 2000K (H2 dissociates). Mass ~ 1 Jupiter; radius ~ several AU, lifetime ~ 102 - 103 yr, ~ 9 candidates.. JWST- some signal at MIR wavelengths… - Class 0: protostar too deeply embedded for source detection, but has outflow.
Dunham+2014
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2. Disk formation
Cloud to YSOs: Specific angular momentum j as a function of scale and YSO evolutionary stage: - j conserved on core infall scales - Keplerian inside 100 AU
Belloche 2013
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Disk formation in turbulent, magnetized cores: 1. Quiescent rotating core with B field: Magnetic braking “catastrophe” NO DISK FORMS! (Mellon & Li 2008, Hennebelle & Fromang 2008, review Z-Y Li+ 2014) - Turbulent reconnection? (Santos-Lima+ 2012,2013) - Turbulent MHD – degrades torque, disk forms – accretion from a few
filaments covering only 10% of area (Seifried+2012, 2013, 2015)
1300 AU scale: 2.6 and 100 M¤, no initial rotation, velocity vectors, B (black lines), forming Keplerian disk (blue), filaments (green). Seifried, Banerjee, Pudritz, & Klessen ( 2015)
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Observational Tests…
MIR: Class 0 observations: Simple envelope profiles ρ ~ r-2 , r-1.5 must be modified (Jørgensen+2005, Enoch+2009)
Complex environments on 1000 AU scales (Tobin+ 2010) - asymmetries and filaments… NOT in line with spherical collapse pictures. JWST MIR observations to identify/image filamentary build up of disks…
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- Upper; magnetic tower flow - Lower; zoomed in by 1000, centrifugally driven disk wind
3. Outflows Early stages: gravitational collapse of rotating, magnetized cores produces disks and disk winds (Banerjee & Pudritz 2006):
2 components of the flow
MHD outflows associated with formation of first hydrostatic core (Bate+2014)
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Solving angular momentum problem for disks and accreting YSOs :
Gressel+2014
Centrifugally driven disk winds carry angular momentum from disk (review Pudritz+2007); Accretion powered stellar winds spin down YSOs (Matt & Pudritz 2005, Zanni & Ferreira 2009)) [disk locking picture – Königl 1991] Disk turbulence (MRI) completely quenched in inner regions by non-ideal MHD (“dead zone”) – ang. momentum transport by disk winds (Bai & Stone 2013, Gressel + 2014..)
Matt & Pudritz (2005)
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How star formation controls planet formation…
Stellar irradiation of evolving disks - temperature structure -> heat transition: viscous to radiative -> position of ice lines Stellar ionizing flux on evolving disks – ionization driven disk chemistry -> position of dead zones All 3 of these act as planet traps – slow planetary migration – affects planetary composition… (eg. Hasegawa & Pudritz 2011, 2012, Alessi, Pudritz, & Cridland 2016, Cridland, Pudritz, & Alessi 2016).
ALMA partnership 2015
Zhang+2015
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Summary:
JWST can drive transformative research on: - Physics of core and filament structure and formation - IMF at BD and Jovian mass scales - HII regions/ UV escape from GMCs - Existence of earliest protostellar hydrostatic states (FHCs) - Filamentary flows and formation of disks - Structure and fragmentation of massive disks - Early outflows