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What Physics Determines the Peak of the IMF? Insights from the Structure of Cores in Radiation-Magnetohydrodynamic Simulations
Astrophysics of Galaxies · Submitted 15 Mar 2016
abstract compiled · 2026-09-23 · no language-model review
The authors' main claim We show that, on small scales, thermal pressure enhanced by stellar radiation heating is the dominant mechanism limiting the ability of the gas to further fragment.
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The abstract, sentence by sentence
As star-forming clouds collapse, the gas within them fragments to ever-smaller masses. Naively one might expect this process to continue down to the smallest mass that is able to radiate away its binding energy on a dynamical timescale, the opacity limit for fragmentation, at $\sim 0.01$ $M_\odot$. However, the observed peak of the initial mass function (IMF) lies a factor of $20-30$ higher in mass, suggesting that some other mechanism halts fragmentation before the opacity limit is reached. In this paper we analyse radiation-magnetohydrodynamic simulations of star cluster formation in typical Milky Way environments in order to determine what physical process limits fragmentation in them. We examine the regions in the vicinity of stars that form in the simulations to determine the amounts of mass that are prevented from fragmenting by thermal and magnetic pressure. We show that, on small scales, thermal pressure enhanced by stellar radiation heating is the dominant mechanism limiting the ability of the gas to further fragment. In the brown dwarf mass regime, $\sim 0.01$ $M_\odot$, the typical object that forms in the simulations is surrounded by gas whose mass is several times its own that is unable to escape or fragment, and instead is likely to accrete. This mechanism explains why $\sim 0.01$ $M_\odot$ objects are rare: unless an outside agent intervenes (e.g., a shock strips away the gas around them), they will grow by accreting the warmed gas around them. In contrast, by the time stars grow to masses of $\sim 0.2$ $M_\odot$, the mass of heated gas is only tens of percent of the central star mass, too small to alter its final mass by a large factor. This naturally explains why the IMF peak is at $\sim 0.2$ $M_\odot$.
background contribution method result interpretation main claim
| # | sentence | role | the authors' own work | states the main conclusion | strength |
|---|---|---|---|---|---|
| s0 | As star-forming clouds collapse, the gas within them fragments to ever-smaller masses. | background 1.00 | 0.05 | 0.01 | asserted 1.00 |
| s1 | Naively one might expect this process to continue down to the smallest mass that is able to radiate away its binding energy on a dynamical timescale, the opacity limit for fragmentation, at $\sim 0.01$ $M_\odot$. | background 1.00 | 0.05 | 0.02 | speculative 0.85 |
| s2 | However, the observed peak of the initial mass function (IMF) lies a factor of $20-30$ higher in mass, suggesting that some other mechanism halts fragmentation before the opacity limit is reached. | gap 0.52 | 0.05 | 0.05 | hedged possibility 0.91 |
| s3 | In this paper we analyse radiation-magnetohydrodynamic simulations of star cluster formation in typical Milky Way environments in order to determine what physical process limits fragmentation in them. | contribution 0.68 | 0.91 | 0.04 | asserted 0.63 |
| s4 | We examine the regions in the vicinity of stars that form in the simulations to determine the amounts of mass that are prevented from fragmenting by thermal and magnetic pressure. | method 0.81 | 0.95 | 0.03 | asserted 0.78 |
| s5 | We show that, on small scales, thermal pressure enhanced by stellar radiation heating is the dominant mechanism limiting the ability of the gas to further fragment. | result 1.00 | 0.95 | 0.85 | asserted 1.00 |
| s6 | In the brown dwarf mass regime, $\sim 0.01$ $M_\odot$, the typical object that forms in the simulations is surrounded by gas whose mass is several times its own that is unable to escape or fragment, and instead is likely to accrete. | result 0.99 | 0.94 | 0.11 | asserted 0.82 |
| s7 | This mechanism explains why $\sim 0.01$ $M_\odot$ objects are rare: unless an outside agent intervenes (e.g., a shock strips away the gas around them), they will grow by accreting the warmed gas around them. | interpretation 0.99 | 0.72 | 0.16 | hedged possibility 0.50 |
| s8 | In contrast, by the time stars grow to masses of $\sim 0.2$ $M_\odot$, the mass of heated gas is only tens of percent of the central star mass, too small to alter its final mass by a large factor. | result 0.95 | 0.91 | 0.09 | asserted 1.00 |
| s9 | This naturally explains why the IMF peak is at $\sim 0.2$ $M_\odot$. | interpretation 0.97 | 0.73 | 0.57 | asserted 0.97 |
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The paper
Original abstract
As star-forming clouds collapse, the gas within them fragments to ever-smaller masses. Naively one might expect this process to continue down to the smallest mass that is able to radiate away its binding energy on a dynamical timescale, the opacity limit for fragmentation, at $\sim 0.01$ $M_\odot$. However, the observed peak of the initial mass function (IMF) lies a factor of $20-30$ higher in mass, suggesting that some other mechanism halts fragmentation before the opacity limit is reached. In this paper we analyse radiation-magnetohydrodynamic simulations of star cluster formation in typical Milky Way environments in order to determine what physical process limits fragmentation in them. We examine the regions in the vicinity of stars that form in the simulations to determine the amounts of mass that are prevented from fragmenting by thermal and magnetic pressure. We show that, on small scales, thermal pressure enhanced by stellar radiation heating is the dominant mechanism limiting the ability of the gas to further fragment. In the brown dwarf mass regime, $\sim 0.01$ $M_\odot$, the typical object that forms in the simulations is surrounded by gas whose mass is several times its own that is unable to escape or fragment, and instead is likely to accrete. This mechanism explains why $\sim 0.01$ $M_\odot$ objects are rare: unless an outside agent intervenes (e.g., a shock strips away the gas around them), they will grow by accreting the warmed gas around them. In contrast, by the time stars grow to masses of $\sim 0.2$ $M_\odot$, the mass of heated gas is only tens of percent of the central star mass, too small to alter its final mass by a large factor. This naturally explains why the IMF peak is at $\sim 0.2$ $M_\odot$.
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How the literature uses it
4 arXiv papers cite this in 16 sentences
Papers citing it · 13 sentences without a confident label
Uses its method or tool 1 paper
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For each shell, we compute the total gas enclosed $m_gas$, excluding the sink particle mass, the mean density $ρ$ of the enclosed material, and its mass-weighted mean sound speed $c_s$.[As in [ref] , we focus on cumulative rather than differential quantities because they are less subject to noise; see Appendix A of [ref] for a demonstration that this choice does not alter the qualitative results.] From the resulting profiles in $ρ$ and $c_s$ we compute an effective temperature [equation] and, more importantly from the standpoint of determining whether gas will fragment, the Bonnor-Ebert mass [ref] , [equation] The Bonnor-Ebert mass characterizes the degree of thermal support.
The Effects of Magnetic Fields and Protostellar Feedback on Low-mass Cluster Formation · Analysis Methodology · 0.97
Compares with it 1 paper
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[ref] noted a similar trend in models of denser high mass star forming regions, and also found that the temperature profiles were always close to $T_eff∝r^-0.3$.
The Effects of Magnetic Fields and Protostellar Feedback on Low-mass Cluster Formation · Profiles of Mean Density, Temperature and Magnetic Field · 0.96
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