Adam Ginsburg's blog

Articles by Adam Ginsburg

• PPVI Talk 12: Clarke

  MULTIPLICITY IN EARLY STELLAR EVOLUTION

  B. Reipurth (University of Hawaii, Institute for Astronomy, Hilo, United States), A.P. Boss (Carnegie Institution, Department of Terrestrial Magnetism, United States), C.J. Clarke (University of Cambridge, Institute of Astronomy, United Kingdom), S.P. Goodwin (University of Sheffield, Department of Physics and Astronomy, United Kingdom), L.F. Rodriguez (Universidad Nacional Autonoma de Mexico, Mexico), K.G. Stassun (Vanderbilt University, Department of Physics and Astronomy, United States), A. Tokovinin (Cerro Tololo Inter-American Observatory, Chile), H. Zinnecker (NASA Ames Research Center, SOFIA Science Center, Germany)

  Observations from optical to centimeter wavelengths have demonstrated that multiple systems of two or more bodies is the norm at all stellar evolutionary stages. Multiple systems are widely agreed to result from the collapse and fragmentation of cloud cores, despite the inhibiting influence of magnetic fields. Surveys of Class 0 protostars with mm interferometers have revealed a very high multiplicity frequency of about 2/3, even though there are observational difficulties in resolving close protobinaries, thus supporting the possibility that all stars could be born in multiple systems. Near-infrared adaptive optics observations of Class I protostars show a lower binary frequency relative to the Class 0 phase, a declining trend that continues through the Class II/III stages to the field population. This loss of companions is a natural consequence of dynamical interplay in small multiple systems, leading to ejection of members. We discuss observational consequences of this dynamical evolution, and its influence on circumstellar disks, and we review the evolution of circumbinary disks and their role in defining binary mass ratios. Special attention is paid to eclipsing PMS binaries, which allow for observational tests of evolutionary models of early stellar evolution. Many stars are born in clusters and small groups, and we discuss how interactions in dense stellar environments can significantly alter the distribution of binary separations through dissolution of wider binaries. The binaries and multiples we find in the field are the survivors of these various destructive processes, and we provide a detailed overview of the multiplicity statistics of the field, which form a boundary condition for all models of binary evolution. Finally we discuss various formation mechanisms for massive binaries, and the origin of massive trapezia and their role in the dynamical evolution of clusters.

  

  Intro

  Most stars are single (M-dwarves), but most were born multiple

  • Many new binary surveys completed
  • simulations produce binary statistics (too) easily
  • new focus on high order multiples, ultra-wide systems

  Why binaries?

  Mass estimation: done using evolutionary tracks

  • means are well-centered
  • scatter around tracks is very wide

  Binaries useful as laboratories

  • Rosenfeld 2012 - 13CO disk + spectroscopic binary
  • mass matching between disk and binary

  Coevality to study differential disk dispersion

  • all reasonable theories say binaries arecoeval
  • disk lifetimes are shorter in closer binaries (Cieza 2009, Harris 2012)
  • Test mass dependence of disk dispersal (Kneller & Clarke 2013)

  New Surveys

  Duquenny & Mayor 1991 was the old standard Now Raghavan is.

  • # of triples doubled
  • Binary fraction increases towards higher masses

  O-stars:

  • close binaries very common
  • Sana 2012: 1/4 of O-stars should merge
  • many others should interact

  M-stars:

  • "early" and "late" M-stars (large dynamic range in masses)
  • smooth trend towards more equal mass ratios towards lower-mass objects
  • separation larger for M than brown dwarf

  Multiplicity fraction: # of stars not single Companion Star Fraction: # of companion stars per system

  • Chen 2013: companion fraction decreases with age from Class 0 (high) to I to main sequence
  • Many caveats on class 0 companions (e.g. Maury et al 2010 - not clear they are multi-stellar systems)

  Class 0 Non-hierarchical triple (Rodriguez, Reipurth 2013)

  • very unstable
  • [what is the line-of-sight separation?]
  • end state is a close binary and an ejection (e.g. Reipurth 2000)

  Class I - Connelly 2008

  • Decline in binary fraction with spectral index
  • probably losing companions as a function of time
  • lose companions at orbital velocity
  • 9 of 9 binaries have a companion near enough to be ejected by this mechanism
  • early days of the field

  Simulations

  Reproducing binary fractions is apparently easy

  • Bate 2009: multiplicity by mass = perfect match * surprising since there's not much physics * Maybe it's just not that sensitive to physics * metallicity, feedback seem not to change binary stats
  • Maybe we're just lucky, though?
  • Gobs of simulations papers cited (Bate, offner, Price, Hennebell, Machida, Burzle, Joos, Seifreid, Boss, Commercon...)
  • Don't ignore magnetic fields even though clouds are supercritical (there's plenty of B-field energy)

  Do B-fields cause problems?

  • Hennebelle & Teyssier 2008: braking so efficient that prevents formation of disk * therefore prevents binary fragmentation * braking drives mass ratio down by reducing angular momentum and accreting onto primary star
  • B-field alignment may be the problem
  • turbulence reduces efficiency of B-braking (see Talk 5)

  Simulations are not well-converged

  What happens if you add MRI?

  • Shi 2012: more vigorous accretion onto binary
  • adds to uncertainty

  Gravitational Dynamics

  Interaction between binary and its environment

  • Are binaries processed by cluster?
  • Can they constrain environment?
  • Need good N-body simulations

  • King 2012 - complete only for Taurus, Cham I, Oph, IC 348, ONC, with range of number densities

    • limited range 62-620 AU
    • no apparent different in multiple fraction (~20%)

  • binaries in clusters maybe excess at small separations... which is not expected?

  Exotica

  Ultrawide tail of Raghavan distribution

  • separations > natal cloud cores
  • Dhital looking for proper motion pairs

  Surprising mode of creating wide binaries in clusters

  • Kouwenhoven 2010, Moeckel & Bate, Moeckel & Clarke
  • temporarily bound pairs many many times
  • if cluster expanding, "transient" bound to neighbor -> actually bound over long timescale

  Statistics of Higher Order Multiples

  Tokovinin et al in prep... (initial results)

  • multiple systems fill all phase space that is stable
  • not consistent with pure N-body decay of initially non-hierarchical systems
  • pure n-body -> high eccentricity, and most often the ejected fellow is small * observationally, this doesn't quite work... * to get away from stability boundary, need gas (Clarke & Pringle 1991)
  • Reipurth & Mikkola 2010, 2012 -> populate binary and triple diagrams
  • PP7 will compare data to these simulations
  • Within triples: * Easier to make high mass ratio binary pairs at early times (1 Myr) * low-mass ratio binary pairs late (100 Myr)

  Questions

  • Q: N-body systems required to form wide binaries... not really, not the only way. In low-density environments, born that way. Even though you can make wide binaries in clusters, it's not the only mechanism. I am concerned with N-body codes - you can create whatever you wish with N-body.
  • A: There are other dynamical routes. You CAN invoke 3-body interaction. Otherwise you need a 10^5 AU core rotating with beta=1
  • Q Kastner: Circumbinary system mislabeled. But, illustrates a ton of interesting things. Hierarchical multiple, surrounded by massive circumstellar disk. 10 Myr (if part of beta pic), far-companion that is also a double.
  • Q Kastner: Gas-rich - how gas rich must it be?
  • A: Orbital evolution, a small amount of mass can go a long way. For formation, obviously must be comparable to mass of star.
  • Q Susanne ?: Binary fraction decreases with age. Why?
  • A: Reconfiguration of unstable multiples. Don't have much time from Class 0 to Class I. 10% in surviving triples, many more form in triples
  • Q Tobin: Binary separations are different in ONC and Taurus. Is that still current or is that an obsolete result?
  • A: No wide proper motion pairs in ONC. Maybe still true?
  • Q: AK Sco - xmm newton, periastron passage -> sudden outburst. See my poster.
  • Q: Mike Myers: Change in q distribution as a function of primary mass? We found q consistent except for BD distribution.
  • A: Not monotonic. O-stars have preference for q=1, so do low-mass stars. If purely dynamic, don't want both.
  • Q MM: in Science paper, they chose flat q....

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• PPVI Talk 11: Matthew Bate

  STAR CLUSTER FORMATION AND FEEDBACK

  M. Krumholz (University of California, Santa Cruz, Astronomy, Santa Cruz, United States), M. Bate (Exeter University, Astrophysics, United Kingdom), H. Arce (Yale University, Astronomy, United States), J. Dale (Czech Academy of Sciences, Mathematics, Physics, and Earth Sciences, Czech Republic), R. Gutermuth (University of Massachusetts, Department of Astronomy, United States), R. Klein (University of California, Berkeley, Astronomy, United States), Z.-Y. Li (University of Virginia, Astronomy, United States), F. Nakamura (National Astronomical Observatory of Japan, Japan), Q. Zhang (Harvard-Smithsonian Center for Astrophysics, United States)

  Stars do not generally form in isolation. Instead, they form in clusters, and in these clustered environments newborn stars can have profound effects on one another and on their parent gas clouds. Feedback from clustered stars is almost certainly responsible for a number of otherwise puzzling facts about star formation: that it is an inefficient process that proceeds slowly when averaged over galactic scales; that it produces mostly unbound field stars rather than bound clusters; and that it produces an IMF with a distinct peak in the range 0.1 - 1 Msun, rather than an IMF dominated by brown dwarfs. In this review we summarize current observational constraints and theoretical models for the complex interplay between clustered star formation and feedback.

  Intro

  • 30 Dor - need for feedback treatment
  • Wrong results without feedback * SFR too high * IMF has too many small stars
  • clusters stay bound unless feedback included. Too bound?

  3 Categories of feedback:

  • momentum: Jets, outflows, radiation pressure. Efficient cooling
  • hot gas feedback: winds from hot stars, high velocity shocks, ionized regions
  • Thermal feedback: heating, thermal pressure

  Momentum Feedback

  Protostellar outflows

  • outflow momentum / cloud mass = velocity; compare to velocity dispersion
  • If all simultaneous, could unbind cloud
  • if gradual, stir turbulence
  • Low coupling could just punch holes

  NGC 1333 "prototypical protocluster"

  • most observations: can drive turbulence
  • outflow power vs turbulence dissipation: often comparable
  • insufficient in GMCs / large scales

  Simulations:

  • outflows can maintain virial equilibrium
  • collimated flows more efficient than spherical flows at driving turbulence
  • jet into smooth medium bad at exciting supersonic turbulence
  • Wang 2010: outflows, turbulence, b-fields all reduce SFR

  Radiation Pressure

  Direct transfer of momentum via photons

  • dominated by massive stars
  • drive winds on galactic scales
  • less important than outflows for clusters with M<10^4 msun

  Feedback significant if \(\dot{p} * t_{ff} \gtrsim M v_{virial}\).

  • \(\dot{p} = f_{trap} L/c\).
  • what is f_trap?
  • Murray says f > 1, others say 1
  • Shane Davis VET vs FLD radiation
  • Efficiency still uncertain

  Hot Gas Feedback

  Hot stars: T>40,000 K

  • M > 40 msun
  • wind momentum flux less than radiation
  • winds create shocks, very hot gas. Cooling is slow
  • But, wind-shocked gas can "leak"
  • HII regions are also important

  Wind observations:

  • Chandra 30 Dor & others
  • Partially leaky
  • HII regions dominate in general
  • ?x-ray dominates in 30 dor? (inconsistent with slide)

  Wind theory:

  • spherically symmetric models
  • leakage parameters
  • simulations: Rogers & Pittard
  • very leaky

  Photoionization / Hot Gas

  Spitzer standard treatment

  Hundreds of HII regions known...

  Theory:

  • limit growth of OB stars
  • Walmsley: accretion flows stall HII region, crush it
  • Keto: gravitational trapping

  HII expansion into clouds

  • disrupts low-mass clouds
  • Dale: GMC survives if escape velocity exceeds 10 km/s

  Triggering

  • weak triggering: same # of stars, but faster
  • strong triggering: increase in overall SFE
  • Top-heavy IMF? No

  RDI: Bisbas, Haworth

  Evidence is weak

  • age gradients: questionable
  • higher sfe: questionable
  • causal links in ages
  • young stars near bubbles
  • 10-20% (Kendrew)
  • negative feedback wins

  SNe feedback

  hot gas feedback

  • not important because delayed by 4 years
  • only 1 sne in 30 Dor, but lots of ejection by other forms
  • very important on large / galactic scales

  Thermal Feedback

  Direct stellar luminosity

  • some nuclear, some accretion
  • <3 msun, accretion luminosity dominates
  • if accretion energy goes into outflows, not thermal
  • accretion bursty would reduce thermal feedback
  • substantial heating in inner portions of clusters
  • thermal feedback inhibits fragmentation: changes IMF
  • reduce dependence on initial conditions for IMF
  • can reproduce IMF

  Cooperative accretion

  • thermal feedback prevents additional fragmentation -> top-heavy IMF
  • everything accretes onto current stars

  more physics in more simulations

  • outflows reduce accretion, which reduces thermal feedback
  • reduces "overheating" problem

  B-fields

  • may increase accretion rates
  • on large scales, inhibit fragmentation
  • overall make star formation efficiency lower

  Questions

  • Q: Chris McKee: Possible role of supernovae. Photoionization occurs first. Supernova goes off in low density cavity, no effect on cloud.
  • A: Potential problem, we just don't know at the moment. In the 10^5+ Msun clouds, the cloud stays intact - HII region doesn't kill it. In those cases, they hit lots of mol. gas
  • Q: Peter ?: Question about methods. Are the sink particles magic? Is this magic?
  • A: Well, we need them in order for the simulations to work. Try to make them as small as possible. Go to first hydrostatic core phase. I think they're not the dominant uncertainties
  • Q: Diederich Kruijssen - How does feedback affect cluster formation? If you end with a bound cluster, feedback didn't matter. Do we understand the degeneracies well enough to make "perfect" simulations.
  • A: No conclusions. This is about where we are now. How does feedback affect the cluster? We still don't know.
  • A2: No, we don't understand the degeneracies between the types of feedback at all.
  • Q: Mike Myer - Global collapse vs driven turbulence. Modest age spreads. SFR too fast & too efficient; I would think either or. Which solves too fast, which solves too efficient?
  • A: Too fast - solved by Wang et al. All feedback & mechanisms, SFR drops a lot. But, none stop star formation in the end. Need some way to disperse cloud still. Need HII regions and other feedback in more massive clusters.
  • Q: Hans Zinnecker - SN feedback. Some evidence in ScoCen for SN triggering because low-mass stars are coeval with other stars. Coeval = must be triggered. Also have to think about starburst galaxies - maybe different mode. Maybe positive triggering in starbursts

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• PPVI Talk 10: Jeffries

  AGES OF YOUNG STARS

  D. Soderblom (Space Telescope Science Institute, Baltimore, United States), L. Hillenbrand (Caltech, Pasadena, United States), R. Jeffries (Astrophysics Group, Keele University, Staffordshire, United Kingdom), E. Mamajek (Dept. of Physics & Astronomy, University of Rochester, United States), T. Naylor (School of Physics, University of Exeter, United Kingdom)

  Determining the sequence of events in the formation of stars and planetary systems and their timescales is essential for understanding these processes, yet establishing ages is fundamentally difficult because we lack direct indicators. In this review we discuss the problem of ages for young stars, specifically for those less than ~ 100 Myr old. We start by establishing a reliable scale scale for young stars using the measurement of the Lithium Depletion Boundary (LDB) in young clusters, a method that involves fairly simple and wellestablished physics and observations. We show that LDB ages are consistent with those from the upper main sequence and the main sequence turn-off if modest core convection is included in the models of higher-mass stars. The LDB method can be used to set limits on the age of individual objects but is primarily applicable to groups of stars. We then review the available methods for age estimation, which include kinematic traceback of young groups, placing stars in HRDs, pulsations and seismology, gravity differences, rotation, activity, and lithium abundance as age indicators. We list the known strengths and weaknesses of each of these. Some of these techniques are useful within certain limits (such as age and mass ranges) and others are entirely problematic. We look at the issue of age spreads within star-forming regions and how well they can be quantified. Finally, we offer suggestions and recommendations to guide interpretation of observations, and we show how the current situation can be improved over the next few years.

  Talk

  Can measure radii, luminosity, but never age directly.

  Ages matter quite a lot: use them to age-date everything else

  Ages during the first Gyr of stellar lifetimes.

  Steps:

  1. Categorise stars
  2. Put them in age-order
  3. Estimate age ratios
  4. Estimate absolute ages?

  What methods can we use?

  • No fundamental methods available
  • Nearly fundamental = lithium depletion boundary, kinematic traceback
  • model-depndent: isochrone fitting, asteroseismology, surface gravity, radii, radiative-convective
  • not useful: metallicity

  Lithium depletion

  • fully convective star drags lithium to core, processes lithium
  • Lithium burns at 1e6 K
  • Why semi-fundamental? Very good agreement at which luminosity Li dep. occurs
  • 5-10% error
  • easy to measure at 6701 AA
  • can't go past 200 Myr
  • models agree on luminosity vs age, but disagree strongly on temperature
  • Measure lithium in whole cluster, look for boundary * 8 measurements: 22-122 Myr old * uncertainty from unresolved binaries

  Below 20 Myr? Kinematic ages, independent of stellar physics

  • traceback ages - trace orbits back to closest location
  • expansion age: position vs. velocity.... hubble expansion
  • flyby age: find when object was closest to group [how does this help?]
  • runaway age: to age-date ejected sources. Assume that ejection occurred at birth (sn ejection)

  Example: Beta pic, trace back to a common smaller volume (40 pc across)

  • problems: some stars don't fit
  • is 40 pc across really a reasonable coeval group?
  • actually, expansion is inconsistent with 12 Myr age using revised Hipparcos

  Runaways: what if they ran away from a different cluster?

  • ONC example - must be absolutely sure you know where runaway came from

  Model-Dependent Methods

  Fitting isochrones to data

  • main-sequence turnoff
  • Hayashi->ZAMS
  • requires distance

  Choice: Fit in HR diagram or CMD?

  • HR: can account for extinction, etc.
  • CMD: Statistics are very badly correlated on HR diagram

  Physical uncertainties increase high-mass stellar lifetimes:

  • Rotation
  • Convective overshooting
  • errors are ~30%
  • low masses unaffected

  ZAMS NGC 6530

  • model including uncertainties and binaries
  • requires age 5.5 Myr
  • statistical uncertainties often dominant

  Models without convective overshooting fail to reproduce Lithium ages

  Pre-main-sequence isochrones

  • best to use HR diagram for differential extinction
  • need spectroscopy to get Teff
  • luminosity spread: binary, variability, age spread, accretion
  • luminosity spreads ubiquitous * are they real? * do they imply real age spread * is there one real age?

  ONC: Cannot explain age spread as observational uncertainty

  • observed age spread 0.4 dex (factor of 3)
  • uncertainties can't explain!
  • HR diagram may be set by Accretion history instead of age of star

  HR scatter shrinks after 10 Myr

  models do not agree on age/teff relation

  observations do not agree with models

  Bell: compare PMS and upper main sequence ages.

  • for some models, there is agreement
  • But, still uncertain: What if high mass stars form significantly after low-mass?
  • Order is OK, but spacing is uncertain by factor ~2

  Asteroseismology

  Kepler can't see "p-modes" because of variability

  delta-scuti instability strip... delta-scuti age more precise than HR fit

  spectroscopic measurements of gravity can get rank-ordered measurements

  R-C gap:

  • distance & extinction independent
  • good to ~15 Myr

  Gyrochronology: good for old stars

  • >100 Myr
  • Slowing period for low-mass stars. Age-dependent "Slope"

  Disk-presence clock?

  • 50% uncertainty
  • dangerous until we understand what disperses disks and when?

  Lithium abundance:

  • extremely strong dependence on efficiency of convection
  • useful relative but not absolute age indicator
  • M and K dwarves have very large EW[Li] change
  • ~factor of 2 age indicator

  Conclusions

  • Lithium depletion boundary model-independent, precise (10%), accuracte (10%)
  • Expensive in terms of telescope time
  • Kinematics: currently problematic
  • Runaways: Origin

  Model-dependent ages:

  • Large statistical uncertainties, modest systematic (main seq)
  • pre-main-sequence: absolute entirely unclear

  R-C gap distance independent Empirical ages

  Final thought: Deuterium depletion boundary should be a <10 Myr age indicator. Requires measuring deuteration in BD atmosphere

  Questions

  • Q: Gunther? from COROT. Deuterium driven: model independence stems from a concensus on the initial conditions. Concensus disappears if you consider colors instead of Hayashi track. What is the zero point?
  • A: Zero point uncertain. Star forgets about initial (cloud) conditions after 20 Myr. Helps keep LDB secure.
  • Q Mamacek: Takeaway - Rescaled ages. Half-life of T-Tauri is 3-4 Myr.
  • Q: What observing proposal for D-burning boundary?
  • A: Deuterated water in atmosphere? CrH->CrD? No idea whether it's even remotely feasible.
  • Q Ian Crossfield: Gyrochronology useless for young stars. Why does it work for old stars?
  • A: The models, and the data, converge
  • Q: Self-consistency of ages for Beta-Pic (isochronal, LiDep, etc. agree). Too pessimistic to disregard kinematics. Need self-consistency between methods.
  • A: self-consistency can lead you along the wrong path. Subjectivity in excluding objects.
  • Q John Tobin: Reggiani work on Orion - subregion, or whole central region?
  • A: Not sure, HST treasury...
  • Q: Multiple bursts?
  • A: Foreground population. Take out low-extinction stars, still looks the same. There is an age spread.
  • Q Neal Evans: Ages of Class I, Class 0, all depend on half-life of 2 Myr for T-Tauri (Class II) stars. 2 Myr is the goal.
  • A: 3 Myr is a more conventional half-life
  • Q: Deuterium depletion. Deuterated methane line in the thermal IR. CRIRES maybe?
  • A: Too hot for methane
  • Q Krumholz: Age spread in Orion. Must have age spread at <1 Myr because crossing time is ~0.4 Myr. Even most efficient SF, things happen in a crossing time, not instantaneously.
  • A: Note the log-scale. 10 Myr is pretty broad
  • Q: Advertise poster....
  • Q: 10 Myr is much longer than the crossing of Orion
  • Krumholz: I disagree

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• PPVI Talk 9: Amelia Stutz

  THE EVOLUTION OF PROTOSTARS: INSIGHTS FROM TEN YEARS OF INFRARED SURVEYS WITH SPITZER AND HERSCHEL

  M. Dunham (Yale University, New Haven, CT, United States), A. Stutz (MPIA, Germany), L. Allen (NOAO, United States), N. Evans (UT Austin, United States), W. Fischer (University of Toledo, United States), S.T. Megeath (University of Toledo, United States), P. Myers (Harvard-Smithsonian Center for Astrophysics, United States), S. Offner (Yale University, United States), C. Poteet (Rensselaer Polytechnic Institute, United States), J. Tobin (National Radio Astronomy Observatory, United States), E. Vorobyov (University of Vienna, Austria)

  Stars and planets form from the gravitational collapse of dense molecular cloud cores. The protostellar phase is the period during which mass accretes from the core, through an accretion disk formed by conservation of angular momentum, and onto a hydrostatically supported protostar. It is the phase during which the initial masses of stars and the initial conditions for planet formation are set, thus understanding how protostars evolve is a crucial ingredient for developing a general understanding of star and planet formation. Identification and characterization of protostars has traditionally been hindered by the embedded nature of these objects. Over the past ten years, new observational capabilities provided by the Spitzer Space Telescope and Herschel Space Observatory have enabled wide-field infrared surveys of entire star-forming clouds with high sensitivities, leading to remarkable progress in our understanding of protostellar evolution. We review several key advances in the field over the past decade, focusing both on the observations themselves and the constraints these large-area surveys are placing on theoretical models of star formation and protostellar evolution. We also emphasize several open questions and debates and outline the future observational and theoretical work necessary to further advance the field.

  Intro

  Core -> First hydrostatic core -> protostar

  Luminosity problem

  • Low-luminosity discrepancy, maybe sensitivity

  Review of Classes: submillimeter vs bolometric luminosity, temperature

  • Classes don't necessarily correspond to Stages because of, e.g., geometry
  • Lifetimes of classes - consistent
  • L_submm / L_bol vs T_bol: L_smm/L_bol is a better envelope diagnostic
  • Monotonic evolution may not be true
  • Deviants from "cartoon picture" are at least 10% of all protostars

  PACS Bright Red Sources

  • 18 reddest sources in Orion
  • 70um-bright, consistent with Stage 0
  • VeLLOs not a unique evolutionary stage
  • Too many VeLLOs: they can't be first hydrostatic cores (implies too long a lifetime)

  How do protostars get their mass?

  • The protostellar luminosity problem: Theoretical accretion luminosity ~10x observed
  • None of the inputs are observationally directly constrained
  • m-dot provides the most lever arm, with order-of-magnitude uncertainty

  Core-regulated vs disk-regulated accretion

  • Can rule out isothermal sphere
  • cannot rule out any other models

  Low-luminosity bursts

  • episode accretion
  • How common are they?

  Fossil evidence:

  • Arce ALMA HH46
  • prompt entrainment by episodic jet
  • 100-year spacing between clumps

  DiGiT double-peaked CO2 ice feature

  • irreversible chemical process
  • observed towards vellos: vellos were warm in the past

  What fraction of final protostellar mass is accreted in bursts?

  What happens in disks?

  • Disk rotation is only direct means for measuring protostellar mass

  The role of environment

  • B335 vs. OMC 2/3
  • Luminosities higher, mean separations lower in more active regions

  Questions

  • Q: Mark Krumholz: Can you use the environmental dependence to distinguish disk physics from environmental = core accretion physics?
  • A: disk physics not environmentally dependent, but feeding through envelopes...
  • Krumholz: not external feeding that drives burstiness. Burstiness shouldn't know about environment
  • Q: Multiplicity?
  • A: Tom & student working on constraining with HST data.
  • Q: How do you check?
  • A: I dunno.
  • A: Neal - ???
  • Q: Structure in HH46. May be formed by... something else?????
  • A: Hard to make link between outflows and accretion. This is one nice example of some evidence but is not conclusive on its own.
  • Q: Lifetime of class I's. Maybe lifetimes have 2-order-of-magnitude variation; can't reference to Class II.
  • A: I agree.
  • Q: Inclination. SED models for diskey stars. Any method that can disentangle geometry from other effects?
  • A: avoid the issue by looking at direct detection of cold emission in envelope
  • Ewine: Centrally concentrated HCO+ emission can separate inclined Class I from Class 0. Best is large samples with ALMA though.
  • Q: IRS spectra from protostars. Amount of silicate lines as signature of episodic accretion.
  • A: Don't have a good answer. Don't necessarily look at those features. Only CO2 ice is clear-cut irreversible process. [missed one question]
  • Q: Ewine - what is the key observational distinction between VeLLOs and true first hydrostatic cores.
  • A: Catch young outflows. ALMA.
  • Q Mike Dunham: First cores - lifetimes. 10^4 years is the strict, absolute upper limit for first hydrostatic cores. 0.06-6 first cores in Perseus. No large-scale surveys for first cores.

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• PPVI Talk 8: Zhi-Yun Li

  THE EARLIEST STAGES OF STAR AND PLANET FORMATION: CORE COLLAPSE, AND THE FORMATION OF DISKS AND OUTFLOWS

  Z.-Y. Li (University of Virginia, Astronomy, Charlottesville, United States), R. Banerjee (Hamburger Sternwarte, Germany), R. Pudritz (McMaster University, Department of Physics and Astronomy, Canada), J. Joergensen (University of Copenhagen, Center for Star and Planet Formation, Denmark), H. Shang (Institute of Astronomy and Astrophysics, Taiwan), R. Krasnopolsky (Institute of Astronomy and Astrophysics, Taiwan), A. Maury (ESO, Garching bei Muenchen, Germany)

  The formation of stars and planets are connected through disks. Our theoretical understanding of disk formation has undergone drastic changes since PPV, and we are on the brink of an ALMA-enabled revolution in disk observation. Disk formation - once thought to be a trivial consequence of the conservation of angular momentum - is far more subtle in magnetized gas. In this case, the rotation can be strongly magnetically braked. Indeed, both analytic arguments and numerical simulations have shown that disk formation is suppressed in the strict ideal MHD limit for the observed level of core magnetization, at least for idealized, laminar, axisymmetric cores with ordered magnetic fields. In the theoretical literature, this "catastrophic" magnetic braking characterizes the situation wherein the angular momentum of an idealized collapsing core is nearly completely removed by magnetic braking close to the central object. We will review what is known about this "magnetic braking catastrophe", possible ways to resolve it, and the current status of early disk observations. Possible resolutions include non-ideal MHD effects (ambipolar diffusion, Ohmic dissipation and Hall effect), magnetic interchange instability in the inner part of protostellar accretion flow, turbulence, misalignment between the magnetic field and rotation axis, and depletion of the slowly rotating envelope by outflow stripping or accretion. We will assess the pros and cons of each of these proposed resolutions and put them into an over-arching physical context. Outflows are also intimately linked to disk formation; they are a natural product of magnetic fields and rotation and are important signposts of star formation. We review new developments on early outflow generation since PPV. The properties of early disks and outflows are a key component of planet formation in its early stages and we will review these major connections.

  Outline

  • Difficulty of forming disks in ideal MHD limit
  • Can non-ideal save disk in 2D?
  • Can 3D effects do it?
  • How can you form rotationally supported disks?

  Dimensionless number: \(\lambda = \frac{(M/\phi)}{1/(2\pi G^{1/2})}\)

  Crutcher: B-fields below critical value

  B-field energy > rotational energy

  Strict ideal MHD: flux always traces stars

  "Magnetic Braking Catastrophe" - flow dynamics near split monopole

  [stopped taking notes early - I didn't know there was any problem associated with disk formation difficulties before this talk...]

  Questions

  • Q: Disagree. Dead zone due to high density. Easy to form disk there.
  • Q: Alyssa? Great that there are sophisticated arguments... I want more effort on predictions of what would be observed
  • A: need to find twist of field?
  • Q: Mac Low - cautionary note. Scale of turbulence very deep.
  • Q: Tobin - Is it suspicious that large, resolved disks have B-field perp to outflow direction? Maybe misaligned fields allow formation of large disks. In other model, do you have enough mass left to form binary stars, etc?
  • A: On core scale, field parallel to outflow
  • Q: Megeath - How long do the simulations run? What is the angular momentum distribution? How does it depend on envelope.
  • Q: Predictions are unclear...

---

• PPVI Talk 7: Chris McKee

  MASSIVE STAR FORMATION

  J. Tan (University of Florida, Astronomy, Gainesville, United States), M. Beltran (INAF-Osservatorio Astrofisico di Arcetri, Italy), P. Caselli (University of Leeds, School of Physics & Astronomy, United Kingdom), F. Fontani (INAF-Osservatorio Astrofisico di Arcetri, Italy), A. Fuente (Observatorio Astronomico Nacional, Spain), M. Krumholz (UC Santa Cruz, Astronomy, United States), C. McKee (UC Berkeley, Depts. of Physics & Astronomy, United States), A. Stolte (University of Bonn, Dept. of Astronomy, Germany)

  The enormous radiative and mechanical luminosities of massive stars impact a vast range of scales and processes, from the reionization of the universe, to the evolution of galaxies, to the regulation of the interstellar medium, to the formation of star clusters, and even to the formation of planets around stars in such clusters. Furthermore, the synthesis and dispersal of heavy elements by massive stars plays a key role in the chemical evolution of the cosmos. Achieving a rigorous theoretical understanding of massive star formation is thus an important goal of contemporary astrophysics. This effort can also be viewed as a major component of the development of a general theory of star formation that seeks to explain the birth of stars of all masses and from all the variety of star-forming environments. Two main classes of theories for massive star formation are under active study, "Core Accretion" and "Competitive Accretion". In Core Accretion, the initial conditions of star formation are self-gravitating, centrally concentrated cores that condense from the surrounding, fragmenting clump environment with a range of masses. They then undergo relatively ordered collapse via a central disk to form a single star or a small-N multiple. In this case, the pre-stellar core mass function has a similar form to the stellar initial mass function. In Competitive Accretion, the material that forms a massive star is drawn more chaotically from a wider region of the clump without passing through a phase of being in a massive, coherent core. In this case, massive star formation must proceed hand in hand with star cluster formation. If stellar densities become very high near the cluster center, then collisions between stars could also be involved in forming the most massive stars. We review recent theoretical and observational progress towards understanding massive star formation, considering a range of observed galactic star-forming environments, physical and chemical processes, comparisons with low and intermediate-mass stars, and connections to star cluster formation.

  Outline

  • Observations
  • Theory

  State of Observations

  Plot 1: the Tan plot

  • Focusing in <10^4 solar mass region
  • IRDCs
  • Falgarone size-linewidth relation
  • Virial parameter, virial relation.
  • Non-equilibrium high-latitude clouds
  • GMCs borderline bound

  PPV: Clumps form clusters, cores form stars

  • Ginsburg 2012: no >10^4 Msun starless clusters
  • Butler & Tan 2012: clump vs core surface density

  How identify starless cores?

  • Chemistry: N2D+ vs N2H+ (Fontani 2011)
  • starless cores have higher deuteration (but not always)
  • Tan 2013: ALMA observations of starless cores
    • use extinction techniques
    • approximate virial equilibrium in cores
    • assumptions: ions line width, pressure equilibrium, etc.
  • H13CO+ cores in Cyg X (Csengeri 2012)
  • Many suggestions of virial equilibrium
  • Li, Kauffman etc. show many clouds have alpha_vir << 1. Explanations:
    • Maybe strong B-fields?
    • Assumed a fixed CO abundance
    • Short time immediately before free-fall collapse

  Accretion onto Cores and Clumps

  • Beltran 2006: NH3 emission and absorption in front of UCHII region * Also CH3CN. * 4e-3 - 1e-2 msun/yr
  • Lopez-Sepulchre 5e-3 msun/yr * outflows indicate 1e-5 msun/yr
  • Wyrowski 2012 clump accretion at a a fraction of free-fall velocity
  • Peretto 2013: 0.027 pc, 2e8 cm^-3 massive core

  Disks around Massive stars

  • unambiguous evidence: VLTI observations (Kraus 2010)
  • Ilee 2013: CO bandhead emission * Consistent with keplerian disk * 25% have disks
  • Protostellar jet from 10 msun star (Carrasco-Gonzalez 2010) * suggests hydromagnetic outflows (collimated) from massive stars * there is synchrotron
  • Strong B-field in clump (Girart 2009)

  Theory

  • protostars contain many jeans masses
  • Turbulence important within core?
  • B-fields maybe less important
  • Radiative feedback important (Kahn 1974, Starrfield 1971) * force of photons can exceed gravity

  Is high-mass accretion a scaled-up version of low-mass star formation?

  Yes: (Turbulent Core)

  • mass of star related to mass of core
  • But: Why not fragment further?
  • Where are protostellar accretion disks?

  No: (Competitive Accretion)

  • tidally modified bondi-hoyle accretion
  • Doesn't work in virialized turbulent medium
  • B-fields reduce accretion

  Radiation pressure still a problem?

  Direct collisions?

  • While forming, or after forming
  • Sana 2012: expect high-mass binaries to merge
  • Moeckel & Bonnell: merger theory
  • but they don't affect observed clusters (Moeckel 2011)

  Radiation suppresses fragmentation

  • column threshold for massive star formation
  • B-fields plus radiation effective at suppressing fragmentation (Commercon 2011)

  G35.2N Gemini / SOFIA observations: Zhang 2013

  Rowan Smith 2013: predicted ALMA observations

  Peters 2010: fragmentation induced starvation Keto 2007: bondi-hoyle accretion

  Radiation escape: beaming, cavities

  Observational questions:

  • Does the IMF mirror the core mass function?
  • Are disks present in high-massprotostars?
  • Protostellar luminosity function for massive stars?
  • Is IMF at high end universal?
  • How do you make very massive >150 msun stars?
  • Can massive stars form in isolation?

  Why is there a difference?

  • initial conditions
  • missing physics

  Questions

  • Q: Doug (Johnstone?) - Core mass function: thermal -> 1 star, turbulent -> ??? stars. Observationally, is the column the important quantity?
  • A: Yes, theoretically. Cores won't necessarily make single star (or system). Massive cores are likely to form high and low-mass stars. Under what conditions will most of the mass go in to one or two stars?
  • Q: Radiation feedback during formation. How does this gel with massive stars forming in close binaries?
  • A: All simulations have terrible resolution. Radiative + magnetic suppresses fragmentation
  • Q Ewine: Radiation feedback... 1 g/cm^2. What scale?
  • A: Up to 1000 AU, on larger scale B-field must be more effective.
  • Q Klessen?: [mic problems] Caution about using virial theorem. Complex density-velocity structure. Surface terms should not be neglected.
  • A: If you look at any object and compare to virial, you need time-dependent and surface terms. Still a good way to assess gravity.
  • Surface term can explain virial parameters...
  • No it can't...
  • Q: Hiro? (Taiwan) - Observations. Not signatures of direct accretion onto stars?
  • A: Most accretion was on to star-disk system. How do you get it onto the star? Yes, that's one of the open difficult questions. Observations within 10 AU of star.
  • Q: Rolf Kuiper - Radiative tail instability. Our simulations were confirmed last year by two groups.... [not a question]
  • A: I agree that flux limited diffusion underestimates force. Zheng also disagrees with you
  • Q: Mike Kuhn (sp?): How does the presence of active formation of low-mass stars prior to massive stars affects the core processes?
  • A: Universality of IMF only true averaged over spacetime. Low-mass stars drive turbulence in massive star-forming region.

---

• PPVI Talk 6: John Bally

  THE MILKY WAY AS A STAR FORMATION ENGINE

  Authors:

  S. Molinari (INAF-IAPS, Rome, Italy), J. Bally (University of Colorado, Boulder, United States), J.-P. Bernard (IRAP, Toulouse, France), S. Glover (Heidelberg University, ITA, Germany), P. Martin (University of Toronto, Canada), T. Moore (Liverpool John Moores University, United Kingdom), A. Noriega-Crespo (IPAC-Caltech, Pasadena, United States), R. Plume (University of Calgary, Canada), L. Testi (ESO-HQ, Garching, Germany), E. Vazquez-Semadeni (UNAM, Morelia, Mexico), A. Zavagno (LAM, Marseille, France)

  We review and discuss the latest observational scenario globally emerging from the latest generation of continuum and spectroscopic global-scale surveys of the mid-plane of the Galaxy from infrared to radio wavelengths, in terms of cloud formation and evolution, star and clump formation rate and efficiency throughout the Galaxy, evidence for density thresholds for the formation of protocluster-forming clumps, as well as the role of large-scale filamentary structures in channeling diffuse ISM into gravitationally bound clumps. The new view of the Milky Way thus assembled will be discussed in the context of competing theoretical frameworks distinguishing between "slow" and "fast" formation, outlining some of the challenges that this body of evidence poses to a variety of theoretical scenarios for large-scale star formation. We will review the most recent methods and evolutionary indicators used to classify the star formation activity in molecular clumps, comparing the different possible derivations of the local SFR and discussing its variation as a function of Galactocentric radius. The order-of-magnitude increase in the number of sources (mostly clumps) that are revealed in state-of-the-art infrared and submillimeter Galaxy-wide surveys make such studies possible with an unprecedented statistical significance. We will also review the present status of our understanding of the Galaxy as a global star formation engine in the context of similar studies in extragalactic systems, evaluating how the data on our own Galaxy compare with the proposed global star formation laws that relate properties of the interstellar medium to the star formation rate.

  Intro

  Link the MW to high redshift.

  "Galactic Ecology"

  • Efficiency "<15%"
  • Feedback -> "superbubbles"
  • "oops"
  • solar circle Z-scale oscillation 80 Myr
  • Epicyclic period 170 Myr [what is this orbital scale?)
  • orbital timescale 220 Myr

  Large scale view:

  • Sco-Cen HI supershell (1973 A&AS)
  • Rho Oph filament swept up by Sco-Cen bubble?
  • Sco-Cen has HI, not H-alpha
  • Alpha-Persei blew a bubble for Cosmologists

  Orion CO

  • Integral shaped filament
  • Velocity jump occurs at cluster
  • Vazquez-Semadeni 2008 colliding cloud resembles ISF

  BGPS -> ATLASGAL -> Herschel

  And now let's start the talk...

  HiGal

  HiGal 360 in OT2

  W3/W4/W5

  • W3 OB cluster
  • W4 cometary cloud
  • W5 (name drop)

  NGC 7538

  • filament full of H2 outflows

  W43

  • 50 Orions
  • 3 Myr old
  • older subgroup to the south

  Goodman's "we're out of the plane"

  • Nearby spiral arms below the "plane"

  W51 - Spitzer

  • massive clusters

  Back to M16, M17

  • M17
  • M17swex
  • red square

  Nessie (Jackson)

  • extended Nessie from Alyssa Goodman

  Carina Nebula

  • eta car
  • Trumpler 16

  NGC 3603

  Filament Formation

  "Zeldovich Pancakes" - tend to collapse along minor axes first Converging flows "shadowing" HII regions -> cometary filament = tail Vazquez-Semadeni colliding flows movie

  Filaments in HiGal

  • Schisano 2013 - Herschel filaments is continuously distributed (no threshold effect)
  • Mass per unit length increases with amt of star formation

  Distances

  • Russeil "Bootstrap Morphology"
  • BeSSeL parallax
  • Distance PDF (Ellsworth-Bowers) * showed example DPDFs * Shirley & Schlingman HCO+/N2H+ velocities * description of the algorithm
  • Bessel points colored by arm

  Battersby 2011 "Maybe I'll skip this"

  • How does dust vary with SF signature

  Pestalozzi face-on view of the galaxy

  • mass-luminosity plot

  Finish with the CMZ

  • 80% of gas at n>10^4 in CMZ
  • Bania's clump - strange kinematics
  • Orbits in a bar
  • Inner Lindblad resonance?
  • compare our CMZ to NGC 1097
  • Dust at l > 0, 24um sources at l < 0

  Molinari Ring

  • Vertical oscillation timescale twice orbit timescale
  • Sgr B2 has high proper motion in the plane
  • Battersby temperature / column maps
  • several 10^7 msun of non-star-forming gas

  Back to CMZ:

  • Free-free emission
  • CND / circumnuclear ring
  • Sgr A* has no influence on the Galaxy
  • star formation starves Sgr A*

  Plug for CCAT, ALMA

  Storm the Castle!

  Questions

  • Q: Mordecai - seen fragmentation starvation in massive star formation as well
  • A: Yep. Also, Phil knows disk around black hole is like a disk around a protostar
  • Q: Philippe - Threshold issue. No threshold in Hi-Gal data.
  • A: missed it, someone walked in front of me
  • A: Molinari - Different physics, different scales?
  • Q: Sergei - Fermi bubble feature. Attributed to star formation. In other galaxies... we've seen one other analog? More than enough energy in a Sgr A* outburst.
  • A: Should ask how we distinguish AGN-driven vs. star-formation-driven flows. AGN should be low-mass, light, relativistic launch. Star formation blows out big HI bubbles (high-mass).
  • Same guy: Cosmologists say AGN can blow up whole galaxies
  • Q Hans: Sequential star formation, triggered SF... are we more informed about triggering now than before?
  • A: Triggering is a great model, but observationally indistinguishable. Only way to know is precision kinematics. Need GAIA. It happens, but can't be proved.

---

• PPVI Talk 5: H.B. Li

  THE LINK BETWEEN MAGNETIC FIELDS AND CLOUD/STAR FORMATION

  Authors:

  H.-B. Li (MPIA, Heidelberg, Germany), A. Goodman (Harvard-Smithsonian Center for Astrophysics, United States), M. Houde (University of Western Ontario, Canada), Z.-Y. Li (University of Virginia, United States), G. Novak (Northwestern University, United States), T. Sridharan (Harvard-Smithsonian Center for Astrophysics, United States)

  The question of whether magnetic fields play an important role in the processes of molecular cloud and star formation has been debated for decades. Recent observations reveal a simple picture that may help illuminate these questions: magnetic fields have a tendency to preserve their orientation at all scales that have been probed - from 100-pc scale inter-cloud media down to sub-pc scale cloud cores. This ordered morphology has implications for the way in which self-gravity and turbulence should interact with magnetic fields: both gravitational contraction and turbulent velocities should be anisotropic, due to the influence of dynamically important fields. Such anisotropy is largely observed. Here we review these recent observations and discuss how they can improve our understanding of cloud/star formation.

  Content

  Visualize field morphology: Turbulent or directed?

  Can't determine importantce of B-fields just by looking at the map

  If kinetic dominates B-field, B-fields will simply trace cloud morphology. Else, they will define it...

  Alfvenic Mach number important...

  B-field orientation and strength remains controversial on all scales.

  0. Observational tools

  • Dust grain alignment
  • polarization of thermal emission A_V > 100mag
  • polarization of reflected starlight A_V < 5 mag
  • Zeeman effect circular poln: line-of-sight component
  • linear line polarization

  M33 SMA observations

  • synchrotron show low-density medium
  • 6 most massive clouds (for stronger signal)
  • offset between CO polarization and spiral arms is double-peaked at 0,90 deg

  Crutcher Zeeman observations

  Turbulence lognormal column PDF.

  • Higher column deviations from B-fields? "turning point" is the gravitational critical density

  Orion:

  • B-fields perpendicular to elongation of "cores" (perpendicular to filaments?)
  • Mean core B-fields are approximately parallel across Orion

  For local clouds, B-field of the "ICM" (inter-core medium?) either antiparallel or parallel

  Clouds flatten: collapse only along b-fields. Nice top illustration

  

  Disk formation:

  • B-braking requires coupling between disk and B-fields
  • Neutral and Ion turbulent dispersion is different because of B-induction
  • (HCO+ systematically wider than HCN)
  • Estimate decoupling scale ~100's au

  Recent:

  • Chapman 2013: CSO B-fields perpendicular to disk
  • Hull 2013: CARMA B-fields uncorrelated with disk angles
  • Tomisaka 2011: Large-scales, perpendicular to disk. Small scales, randomized.
    • transition from poloidal to toroidal B-field

  Conclusions

  Orion has ordered B-fields from large to small scales

  In general, large-scale fields lead to squishing along B-fields, small-scale

  

  Questions

  • Q: Benetti - Are you only picking out the large-scale components when there may be stronger small-scale fields? How can you be sure you're seeing the dominant component?
  • A:
  • Q: Strong reduction of degree of poln in core. Are you sampling field in core?
  • A: Poln correlates with line widths?
  • Q Tan: Nice results. Are clouds are sub- or super-critical? If you account for extra turbulence, they are just about critical?
  • Q Tan: What about sampling IMF from field strengths that produce massive cores?
  • A: No idea
  • Q Chris McKee: Padoan's group - let weak B-field amplify in simulations. Only by assuming input strong-B-field do we get correlation between large and small scales.
  • A:

---

• PPVI Talk 4: Christoph Federrath

  THE STAR FORMATION RATE OF MOLECULAR CLOUDS

  Authors:

  P. Padoan (ICREA & Institute of Cosmos Sciences, University of Barcelona, Spain), C. Federrath (School of Mathematical Sciences, Monash University, Australia), G. Chabrier (Ecole Normale Superieure de Lyon, France), N. Evans (The University of Texas at Austin, United States), D. Johnstone (Herzberg Institute of Astrophysics, National Research Council of Canada), J. Jorgensen (Star and Planet Formation Center, University of Copenhagen, Denmark), C. McKee (University of California, Berkeley, United States), A. Nordlund (Niels Bohr Institute & Star and Planet Formation Center, University of Copenhagen)

  We review recent advances in the analytical and numerical modeling of the star formation rate in molecular clouds and discuss the available observational constraints. We focus on molecular clouds as the fundamental star formation sites, rather than on the larger-scale processes that form the clouds and set their properties. Molecular clouds are shaped into a complex filamentary structure by supersonic turbulence, with only a small fraction of the cloud mass channeled into collapsing protostars over a free-fall time of the system. In recent years, the physics of supersonic turbulence has been widely explored with computer simulations, leading to statistical models of this fragmentation process, and to the prediction of the star formation rate as a function of fundamental physical parameters of molecular clouds, such as the virial parameter, the rms Mach number, the compressive fraction of the turbulence forcing, and the ratio of magnetic to gas pressure. Infrared space telescopes, as well as ground-based observatories have provided unprecedented probes of the filamentary structure of molecular clouds and the location of forming stars within them. We may thus be on the verge of a breakthrough in our understanding of star formation, and particularly of the star formation rate of molecular clouds.

  Content

  Star Formation Rate from statistics of turbulence

  Drastic differences in SFR, not just because of column density. (Lada 2010: Pipe vs Rho Oph, Rho = 15 Pipe)

  Intro slides:

  • KS relation with local clouds

  • CMF -> IMF (30-70% SFE)

  • Radiation feedback:

    • only affects small scales
    • affects > 10^8 cm^-3

  • B-fields

    • reduce fragmentation more effectively than radiation
    • drive outflows

  • Outflows & Jets

    • direct: reduce star mass
    • indirect - drive turbulence
    • triggering SF

  What drives turbulence?

  • Jets contribute to turbulence on small scales, but don't feed back effectively on large scales
  • Ionization fronts, bubbles?
  • SN explosions
  • MRI / shear?

  Turbulence

  • Reynolds > 1000

  • cites da Vinci

  • Supersonic "cascades" to sonic scale

    • Kolmogorov -5/3 turbulence is incompressible (subsonic)
    • Larson -2 supersonic, shock-dominated

  Turbulence in a Box

  • Mach numbers 2-50
  • "Forcing term" f
    • divergence free vs. compressive turbulence
    • compressive -> bigger voids, higher overdensity
  • Density PDF: lognormal
  • Perseus observations: similar Mach numbers, different dispersions: requires changed b
  • Column vs. Volume density: Brunt
  • Kainulainen, starless -> taurus -> IRDCs
    • starless: b=0
    • starry: b=0.5
    • IRDCs: higher

  Density PDF -> SFR

  All recent theories of SF, SFR are based on an integral over the density PDF.

  • Hennebelle & Chabrier theory: Different free-fall times at different densities, so must be in an integral

  • Cricitical density \(s_{crit}\propto \ln \left(\alpha_{vir}\mathcal{M}^2\right)\)

  • SFR depends on alpha_vir, b, and M

  • Dependence on alpha_vir: inversely proportional to SFR, SFR propto alpha_vir^-1/2

  • Dependence on Mach number

    • stronger when you consider non-constant t_ff
    • SFR larger as gas compressed more

  • Dependence on b: Compressive is much more star forming

    • compressive drives gas to self-gravitating densities far more rapidly

  • What about magnetic fields? Add beta

    • B-fields "wash out" structures
    • Reduce fragmentation

  • Direct measurement of line-of-sight star counts and gas density

    • 1%, 10% SFE examples
    • makes a distribution similar to Heiderman plot
    • agreement assumes total core-to-star efficiency = 100%

  Conclusions

  • supersonic, magnetized turbulence is important for driving SFR

  Questions

  • Q: Adam Frank - Abstract vs realistic forcing from, e.g., outflows. Driving sources makes turbulence look different.
  • A: These experiments are essential. Need to find out.
  • Q: Fabian Heitsch - How does non-isothermality affect theory?
  • A: We are looking only at density ranges where gas is isothermal. SFR is not very dependent on T, but IMF is.
  • Q: How do you explain that solenoidal is a better fit, when all drivers are compressive?
  • A: Don't know. BUT, compression drives solenoidal modes [this is possibly interesting - does a single compression event drive a lot of solenoids?]
  • Q: What are the effects of field-slippage (non-ideal MHD, ambipolar diffusion)
  • A: Not that important....
  • Q: Tan - Filaments have ordered B-fields on large scales. Do you see that in the simulations
  • A: Fields are very intermittent, bouncing all over the place. Large scale "input" B-fields may be more important for mol. clds
  • Q: Solenoidal -> toroidal fields, which squeeze gas. Do you see squeezing (compression) from solenoidal?
  • A: "simple" version of B-fields only works for magnetic pressure, not magnetic tension. For subsonic, SFR doesn't work. Need disordered fields
  • Q: Clark? - How do these results compare with Lada's?
  • A: Don't know, didn't want the complexity

---

• PPVI Talk 3: Paul Clark

  THE ORIGIN AND UNIVERSALITY OF THE INITIAL MASS FUNCTION

  Authors:

  S. Offner (Yale University, United States), P. Clark (Heidelberg University, ITA, Germany), P. Hennebelle (Observatoire de Paris, LERMA, Paris, France), N. Bastian (Liverpool John Moores University, United Kingdom), M. Bate (University of Exeter, United Kingdom), P. Hopkins (California Institute of Technology, Department of Astronomy, United States), E. Moraux (Institut de Planetologie et d'Astrophysique de Grenoble (IPAG), France), A. Whitworth (Cardiff University, Physics & Astronomy, United Kingdom)

  Abstract

  We review current theories for the origin of the Stellar Initial Mass Function (IMF) with particular focus on the extent to which the IMF can be considered universal across various environments. To place the issue in an observational context, we summarize the techniques used to determine the IMF for different stellar populations, the uncertainties affecting the results, and the evidence for systematic departures from universality under extreme circumstances. We next consider theories for the formation of prestellar cores by turbulent fragmentation and the possible impact of various thermal, hydrodynamic and magneto-hydrodynamic instabilities. We address the conversion of prestellar cores into stars, and evaluate the roles played by different processes: competitive accretion, dynamical fragmentation, ejection and starvation, filament fragmentation and filamentary accretion flows, disc formation and fragmentation, critical scales imposed by thermodynamics, and magnetic braking. We present explanations for the characteristic shapes of the Present-Day Prestellar Core Mass Function and the IMF and consider what significance can be attached to their apparent similarity. Substantial computational advances have occurred in recent years, and we review the numerical simulations that have been performed to predict the IMF directly and discuss the influence of dynamics, time-dependent phenomena, and initial conditions.

  Last time, IMF discussion was very heated. (technical difficulties)

  Intro to IMF

  Salpeter, Kroupa, Chabrier, de Marchi ("Tapered" power law) (aside) Thies & Kroupa: BDs not part of IMF

  SFR indicators only see >10 Msun sources. Must extrapolate downward. This is why extragalactic people care.

  Are there variations in the IMF?

  Locally, no.

  More evolved clusters, yes on the low-mass end:

  • Pleiades, NGC 2516, etc.
  • Lose low-mass stars with age dynamically
  • fraction of BDs remaining drops to zero over time

  Difficulties in measurement:

  • Patchy extinction
  • young clusters not done sampling IMF
  • don't understand hayashi tracks that well

  Modal plot from van Dokkum

  • Hints at change in IMF

  Extragalactic - big camps argue top, bottom heavy and universal

  • rely on mass to light ratios
  • M/L models depend on population synthesis models

  What determines the IMF?

  Hoyle, Larson: the cloud fragmentation process. Clouds fragment into cores, cores make stars.

  Does supersonic turbulence regulate? (movie from Federrath)

  Padoan:

  • CMF driven only by turbulence.
  • # of BDs controlled by Mach number.
  • Cloud-in-cloud problem: small cores may be in big cores

  Press-Schechter & Excursion Set formalism

  • Gravitational decoupling sets "critical" scale

  • Hopkins: Density as a function of scale

    • large scale: shear support
    • intermediate: turbulence
    • small scale: hrm...
    • last-crossing => IMF

  • Hennebelle & Chabrier 2008 - account for non-thermal support

  • GMCs are a separate population

  • Useful: turbulence -> more BDs, so extragalactic maybe more turbulent?

  Does the CMF map onto the IMF?

  Cores must:

  • neither accrete nor merge (or you break the CMF my moving cores)
  • fragment self-similarly: same number of fragments per core
  • Same efficiency regardless of mass
  • formation must happen on similar timescale (can't have high-mass first)
  • Cores must be able to form stars, i.e. must be grav. bound

  Can cores form stars?

  • Dust cores depend on column & temperature & line of sight effects
  • number of Jeans masses goes as T^-3
  • Velocity structure? B-fields?

  Fragmentation within cores

  • Used to think that everything fragmented...
  • strong B-fields reduce fragmentation
  • Still haven't studied parameter space well

  SFE within cores

  • Binaries or triples form, requires 30%-50% efficiency
  • can work with outflows
  • but not clear... may depend on well-aligned B-field

  Problems with CMF:

  • Pipe cores follow CMF shape, but are pressure bound
  • CMF in Aquila is too low
  • need CMF of Class 0 sources

  Clusters

  What effects have cluster simulations had on IMF studies...

  • Initial Conditions imprints thermal jeans mass in box, which imprints onto IMF
  • thermodynamics matter

  Fraction of cloud doesn't matter

  • IMF is same in small periodic box and isolated cloud
  • colliding flows can build clouds, be more consistent
  • How we drive turbulence doesn't seem to affect IMF (solenoidal, compressive same)
  • power spectrum of turbulence doesn't affect IMF

  Metallicity?

  • Myers 2011: 20x metallicity change -> same IMF
  • Multiplicity may change (Bate)
  • even 10^-5 Z_sun -> same IMF

  Things that DO affect IMF:

  • radiative transfer: reduces fragmentation
  • need wind feedback, otherwise radiation -> flat IMF by preventing fragmentation but keeping gas
  • B-fields reduce characteristic mass by creating more filaments with different jeans masses

  State of Simulations:

  • Top of IMF about right
  • bottom is wrong

  Secondary Metrics

  Multiplicity

  • Krumholz & Bate both reproduce observations fairly well
  • Require radiative transfer

  Summary

  • IMF studies focused on low-mass end instead of high
  • Cores better understood

  Future: Put simulations into observational plane

  Questions

  • Q: Properties of turbulence in simulations give same IMF...
  • A: Outflows drive turbulence, but turbulence doesn't change IMF. Outflows remove gas.
  • Q Zinnecker: Binaries.
  • A: Cathy will talk about this. Fragmentation knobs are the issue Unclear that binary properties will stay constant (does field match cluster?) N-body evolution matters
  • Q (Galactic center star studier...): not done with discussion of high-mass issues. Radiatitve transfer still matters.
  • A: Many effects are actually from mass segregation
  • Q: Alyssa Goodman: Are the CMF and IMF both lognormals?
  • A: Sometimes fitting distributions they expect. Maybe different processes in low mass clouds.

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