DTA Meeting (BISTRO-2)

Magnetically Controlled Star Formationin the Galactic Disk
Shu-ichiro Inutsuka (Nagoya University)

24 November 2016

National Astronomical Observatory, Mitaka, Tokyo

M51 Synchrotron

Polarization(Fletcher & Beck 2005)

Various Energy Densities
NGC6946

Beck 2004

equipartition assumed

Star Formation is Inefficient!
• Typical Density observed by CO: nCO ~ 102 /cm3
• Mass of Molecular Clouds (~10K): MMC ~ 109 M
• Free-Fall Time for 102 /cm3 ~ 106 yr• Star Formation Rate (if at Free-Fall Rate)
RSF = 109 M /106yr = 103 M /yrWrong!

Observed Star Formation Rate ~ 100 M /yr
 Either Slow or Very inefficient (~10-3)!

Schmidt-Kennicutt Law of SF
ΣSFR [M /kpc2 yr]

Star Formation Rate

High Density Tracer

Lada+2010,2012,2013

Kennicutt 1998

Σgas [M/pc2]
Timescale: Σgas / ΣSFR ~ Gyr• Column Density: Σgas [M/pc2]• SF Rate: ΣSFR [M /kpc2 yr]• Timescale: Μ /(SFR) ~ 20Myr
See also Gao & Solomon 2004; Wu+2005; Bigiel etal. 2008,2010,2011…

Outline• Formation of Molecular Clouds– Phase Transition Dynamics– Thermal Instability, Sustained Turbulence– Effect of Magnetic Field• Dynamics of Filaments– Mass Function of Dense Cores  IMF• Galactic Picture of Cloud/Star Formation– Accelerated Star Formation– SF Efficiency & Schmidt-Kennicutt Law– Mass Function of Molecular Clouds• Possible Mission of BISTRO

Formation ofMolecular Clouds

Radiative Equilibrium for a given density
Solid: NH=1019cm-2,

Dashed: 1020cm-2

WarmMedium

Cold Neutral
Mediume.g., Wolfire et al. 1995, Koyama & SI 2000

20km/s Shock Propagation into WNM

No forcing! Koyama & SI (2002) ApJ 564, L97

Shock Propagation into WNM
122345Ambient ISM
WNM

High PressureMedium

Color: Density
Koyama & Inutsuka (2002) ApJ 564, L97

Summary of TI-Driven Turbulence• 2D/3D Calculation of Propagation of ShockWave into WNM via Thermal Instability
 fragmentation of cold layer into cold
clumps with long-sustained supersonicvelocity dispersion (~ km/s)
1D:Shock ⇒ Eth ⇒ Erad
2D&3D: Shock ⇒ Eth ⇒ Erad + Ekinδv ~ a few km/s < CS,WNM =10km/s
104K due to Lyα line: Universality!
TCNM ~102K C+158µm (~92K)
Koyama & SI (2002) ApJ 564, L97

xelsc.f.
Kritsuk &Norman1999

Hennebelle & Audit 07

20 pc

10,0002

Heitsch et al.
40962
2DVazquez-Semadeni
et al. 2011

Property of “Turbulence”…SubsonicE(k) ∝ k -5/3

δv < CS,WNM 
Kolmogorov Spectrum
2D: Hennebelle & Audit 2007;
See, e.g., Gazol & Kim 2010

Spectrum Observed in ISM
Chepurnov & Lazarian 2010
Armstrong et al. 1995

Hennebelle & Audit 07

20 pc

10,0002

Magnetic Fieldchanges Story!

c.f.
Kritsuk &Norman 1999

Heitsch+ 2006
2D, 40962
Vazquez-Semadeni
et al. 2011Colliding WNM with B0=3µG
v=10km/s(a) 15deg

10km/s

10km/s

10km/s

10km/s<δB2>init = B02(a) 40 deg<δB2>init = 4B02
Ambipolar Diffusion Included

2-Fluid MHD Simulation (AD included)

Inoue & SI (2008) ApJ 687, 303

Compression of Magnetized WNM
Can direct compression of magnetized WNM
create molecular clouds?  Not at once!
Inoue & SI (2008) ApJ 687, 303
Inoue & SI (2009) ApJ 704, 161
Essentially same result byHeitsch+2009; Körtgen & Banerjee 2015;Valdivia+2016
We need multiple episodes of compression.

Compression of CNM (HI)  H2
Compression along
Magnetic Fieldlines, + H2,CO
Formation ofMagnetized
Molecular Clouds
Transformation of HI to H2
Inoue & SI (2012) 759, 35

Further Compress. of Mole. Clouds
Multiple
Compressions ofMolecular CloudMagnetized
Massive Filaments& Striations

Agree withObservations!

Black Lines: Magnetic Field Lines
Self-Gravity Included, SI, Inoue, Iwasaki, & Hosokawa 2015

Observed Molecular Clouds
Cox, Arzoumanian, André+2016

JCMT-BISTRO???

Yellow and Blue Lines: Magnetic Field Lines

Highlight of Herschel Result (André+2010)

2Cs2/G

Self-Gravity Essential in Filaments

Dynamical Timescales of Star Formation
Observational Demography of YSOs (e.g., Fuller&Myers1985)
 NTTauri / Nprotostar ~101.5-2  Tprotostar ~105 yr
 # of Dense Cores: N+IR / NnoIR ~ 101  Tcore ~106 yrc.f. Tambipolar ~107 yr & Tfreefall ~105 yr for n=104/cc

Gravitational collapse of a core is not quasi-steady!
Dynamical Gravitational Collapse in Dense Cores!
Dynamical Evolution in Self-Gravitating Filament withMline ~ 2Cs2/G

Star Formation Efficiency
Herschel Observation (e.g., Andre+2014)

Mcore / Mfilament <~ 15%Star Formation Efficiency in Dense Core: εcoreεcore ~ 33% (ex. Machida+)
εdense gas= Mcore / Mfilament × εcore ~ 5%Star Formation Efficiency of Dense Gas: εdense gas(106 yr)−1 × εdense gas = (20Myr)−1
 tdense gas ~ 20Myr

Mass Function of Dense Core?
Aquila CMF from HerschelAndré+2010; Könyves+2010

SI & Miyama 1997

Larger Wavelength
 Massive Core

Mass Function of Cores in a Filament
Inutsuka 2001, ApJ 559, L149

Line-Mass Fluctuation of Filaments
Initial Power SpectrumP(k) ∝ k –1.5
Mass Function

dN/dM∝M -2.5
Observation of Both Perturbation
Spectrum and Mass Function
direct test !
(cf. Hennebelle & Chabrier 2008;
Shadmehri & Elmegreen 2011)≈ 5/3: Kolmogorov!

SI & Miyama 1997

P (k) ∝ k -1.5
SI & Miyama 1997

t/tff = 0 (dotted) , 2, 4, 6, 8, 10 (solid)

Obs P (k) ∝ k -1.6 (André+2014 PPVI; Roy+2015)

“A possible link between the power spectrum ofinterstellar filaments and the origin of theprestellar core mass function”
Roy, André, Arzoumanian et al. (2015) A&A 584, A111

δ ...
GaussianP (k)∝kn
n= −1.6±0.3
Supporting Inutsuka 2001

Applicability of Filament Paradigm for Massive Stars?
Aquila CMF from HerschelAndré+2010; Könyves+2010
Massive stars can beformed in filaments?

SI & Miyama 1997

Larger Wavelength
 Massive Core

Massive Stars through Filaments(Peretto+2013)• Uniform but Different Velocity in Each Filament• Infall through Filament ~ 10-3 M/yr
Nicely Understood in Filament Paradigm

Filament Paradigm
Completely Successful?!
?

Other Modes of StarFormation?
Cloud Collision (Fukui, Tan, Tasker, Dobbs,...)
Collect & Collapse (Elmegreen-Lada, Whitworth,Palouš, Deharveng, Zavagno,…)Toward Global Pictureof Star Formation
Multiple Compressions Neededfor Molecular Cloud Formation

tform = a few107yr

Network of Expanding Shells
Multiple Episodes of Compression Formation of Magnetized Molecular Clouds
GMCCollisionDenseHI Shell

Molecular
CloudLong (>10Myr) Exposure Picture!
Each bubble disappears quickly (<Myr).
SI+2015; cf. Elmegreen 2007

Velocity Dispersion of Clouds
Multiple Episodes ofCompression Formation of MagnetizedMolecular Clouds

Shell Expansion
Velocities ~ 101 km/s

Observation
Stark & Brand 1989

Cloud-to-Cloud
Velocity Dispersion

Network of Expanding Shells
Multiple Episodes of Compression 
Formation of Magnetized Molecular Clouds
GMC
Collision

Peretto+2013
Inoue & Fukui 2013
Fukui+2012

DenseHI Shell

MolecularCloud

Each Bubble Visible Only for Short Time (~1Myr)!
δv of Mole Clouds ~ vexp of Shells ~ 10km/s

Accelerated Star Formation

NumberPalla & Stahler 2000

Molecular Cloud GrowthCollisions of Clouds Accelerated SF

Age t (Myr)

Also in Lupus, Chamaeleon,
ρ ophiuchi, Upper Scorpius,
IC 348, and NGC 2264
c.f., Vazquez-Semadeni+2007

Possible Mission of JCMT-BISTRO
Aquila CMF from Herschel

BISTRO Obs

filament ⊥ B?
filament // filament?

 Intermediate MassStar Formation

Könyves+2010

Peretto+2013
Massive Star Formation

Massive Star Formation in Ridge

Battersby+2014

Extensive Herschel Studies on Massive Star
Formation in “Ridges”

Ridge orEdge-On
Shell?

Battersby+2014

Edge-On View ofCompressed Shell
 Ridge or Bar!

Bubbles (cyan dashed circles)
HII regions (cyan solid circles)
SNR 3C 391 (yellow oval)
Wolf–Rayet star WR 121b (red oval)

Advent of Large Surveys such as FUGIN

Numerous Straight
Ridges or Bars! Why?

Edge-On View of Compressed Shells =Ridges or Bars!  Bar // B Obs Proof of Cloud Formation Theory!!!

Summary• Fragmentation of FilamentsCore Mass FunctionIMF• Bubble-Dominated Formation of Molecular CloudsUnified Picture of Star Formation δvcloud-cloud ~ 101km/s
 Accelerated Star Formation
 Star Formation Efficiency: εSF~10-2
 Schmidt-Kennicutt Law
 Slope of Cloud Mass Func =1+𝑇𝑇form /𝑇𝑇dis ~1.7

 Intermediate Mass Star Formation  BISTRO
 Colliding Bubbles  Ridges/Bars  FUGIN
SI, Inoue, Iwasaki, & Hosokawa 2015, A&A 580, A49
Kobayashi+2016 submitted

Destruction ofMolecular Clouds
How to Stop
Star Formation?
Radiative Feedback
Photodissociation Critical!
c.f. Dale, Walch,…

Expanding HII Region in Magnetized
Molecular Cloud

Sh104
Deharveng et al. 2003

Radiation Magnetohydrodynamics Calculation
UV/FUV + H2 + CO Chemistry (Hosokawa & SI 2005, 2006ab, 2007)

Disruption ofMagnetized
MolecularClouds
Feedback due to UV/FUVin a Magnetized Cloudby MHD version ofHosokawa & SI (2005,2006ab)30M star destroys
105M H2 gas
in 4Myrs!

Mg(M*)
Non-Star Forming Gas−β+1=1.3−βCentral+1=1.5Stellar Mass, M* / M
−β+1=1.7
Exponent of IMF

M*−β+1 Mg(M*)
Central Stellar Mass, M* / M(SI, Inoue, Iwasaki, & Hosokawa 2015 A&A 580, A49)

Star Formation Efficiency &Schmidt-Kennicutt-Law
105M molecular cloud destroyed by M* > 30M in 4Myrs!

Suppose Mtotal ∼103M stars formed in 105M ∼1 massive (>30M) star for std IMF103M
εSF = 5= 0.01
10 M

Zuckerman & Evans 1974

Star Formation Time

Cloud Disruption Time: Τd=4Myr+T*
Gas Dissipation time: τdis =TdεSF∼ 1.4GyrNo Dependence
on Mass SchmidtKennicutt Law

Star Formation Efficiency, KS-Law
Mg molecular gas (H2) dispersed by Md*
β: exponent of IMF
M*m : Effective Minimum Stellar Mass

If Mg =105, Md*=30M, M*m=0.1M, β =2.5,103M
εSF = 5= 0.01 observation β of IMF
10 M

(SI, Inoue, Iwasaki, & Hosokawa 2015 A&A 580, A49)

Galactic Population ofMolecular Clouds

Mass Function of Molecular Clouds
𝑑𝑑𝑑𝑑 = 𝑁𝑁cl 𝑀𝑀cl 𝑑𝑑𝑀𝑀cl
𝑁𝑁cl
𝜕𝜕𝑁𝑁cl𝜕𝜕
𝑑𝑑𝑀𝑀cl
+𝑁𝑁cl
= 0−
𝜕𝜕𝑡𝑡
𝑑𝑑𝑡𝑡
𝜕𝜕𝑀𝑀cl𝑇𝑇dis
In steady state→ 𝑁𝑁cl 𝑀𝑀cl

𝑁𝑁0 𝑀𝑀cl=
𝑀𝑀0 𝑀𝑀0

𝑀𝑀cl
𝑇𝑇form−𝛼𝛼𝑇𝑇dis =const.
“KS Law”

𝑇𝑇form, 𝛼𝛼 = 1 +𝑇𝑇dis

𝑇𝑇dis ~14Myr & 𝑇𝑇form ~10Myr → 𝛼𝛼 = 1.7(SI, Inoue, Iwasaki, & Hosokawa 2015 A&A 580, A49; Kobayashi+2016 submitted)Summary• Fragmentation of FilamentsCore Mass FunctionIMF• Bubble-Dominated Formation of Molecular CloudsUnified Picture of Star Formationδvcloud-cloud ~ 101km/s
Accelerated Star Formation
Star Formation Efficiency: εSF~10-2
Schmidt-Kennicutt Law
Slope of Cloud Mass Func =1+𝑇𝑇form /𝑇𝑇dis ~1.7

SI, Inoue, Iwasaki, & Hosokawa 2015, A&A 580, A49
Kobayashi+2016 submitted

How far can we trace back?
Redshift011

gas fraction

10

Solar System5Present0.2

SFR

14 13 12 11 10 90.4 0.6 0.8
gas fraction

200SFR (Msun/yr)

1538765Age [Gyr]43210Haywood+2016

Present SF mode responsible for z < 2in Our Galaxy!

©Annie Hughes, MPIA

©Annie Hughes, MPIA

Slope of Cloud Mass Function
Steady State Mass Function of Molecular Clouds→ 𝑁𝑁cl 𝑀𝑀cl

𝑁𝑁0 𝑀𝑀cl=
𝑀𝑀0 𝑀𝑀0−𝛼𝛼𝑇𝑇form, 𝛼𝛼 = 1 +𝑇𝑇dis

Typically, 𝑇𝑇dis ~𝑇𝑇form + 4Myr → 𝛼𝛼 = 1.7
In low density region (Inter-Arm Region)
Larger Tform > Tdis  Larger αIn high density region (Arm Region)
Smaller Τform  Smaller α GMCs in M51 (Colombo+2014)

©Annie Hughes, MPIA

Characteristic Widths of Filaments

Which is determinant, NH or Filament-Width?

AquilaPolarisHerschel filaments have almost the same radii!
Aquila: 2R=0.1pc & ML = 2Cs2/G  NH ≈ 1022cm-2 (Av= several)
Polaris: 2R=0.1pc & ML < 2Cs2/G  NH < 1022cm-2 (Av< several)“Column Density Threshold” is a consequence?