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Chapter 6 Population III Star Formation Under the Impact of Cosmic Rays 1 6.1 Overview SN explosions are thought to be the site of cosmic ray (CR) production (e.g. Ginzburg and Syrovatskii 1969), and the CR background built up from the first SNe may affect the cooling and collapse of primordial gas clouds. Earlier studies (Shchekinov and Vasiliev 2004; Vasiliev and Shchekinov 2006) have shown how the presence of CRs at high redshifts could have lowered the minimum mass at which primordial gas could cool and collapse into virialized dark matter (DM) haloes (e.g. Haiman et al. 1996; Tegmark et al. 1997). This earlier study assumes the existence of ultra-heavy X particles that decay into ultra-high-energy CRs which interact with the cosmic microwave background (CMB), leading to the production of ionizing photons. In this paper, however, we instead examine CR effects which occur through direct collisional ionization of neutral hydrogen. The free electrons created from the ionization act as a catalyst for the formation of H2 (McDowell 1961). Because molecular hydrogen emits photons through its rovibrational transitions, H2 is able to cool primordial gas at temperatures below the threshold for atomic hydrogen cooling (< ∼ 10 4 K). Furthermore, CR ionization can indirectly lead to higher 1 Large portions of this chapter have been previously published as Stacy A. & Bromm V., 2007, MNRAS, 382, 229. Reproduced by permission of John Wiley and Sons. 144
abundances of HD, which also acts as a cooling agent in low-temperature primordial gas (see Johnson and Bromm 2006). However, the presence of CRs can lead to ionization heating, as well. Whether this direct heating effect is strong enough to counter the additional cooling must be determined and will depend on the high-redshift CR energy density. Similar to our work, Rollinde et al. (2005, 2006) used models of early star formation to estimate the CR energy density in the early Universe, constraining it with the observed Li abundances in metal-poor Galactic halo stars. In this paper, we will investigate the importance of the CR feedback on Pop III star formation by modeling its effect on the cooling of primordial gas in two cases: collapse within minihaloes, and shocks associated with the virialization of more massive DM haloes during the later stages of structure formation. The outline for this paper is as follows. In Chapter 6.2 we discuss CR acceleration and propagation in the high-redshift Universe and how these might differ from the present-day case. Chapter 6.3 describes the evolution of primordial gas in minihaloes and in virialization shocks when accounting for the effects of CRs. For the minihalo case we furthermore discuss how the fragmentation scale could change for a sufficiently high CR flux. We present our conclusions in Chapter 6.4. 6.2 Cosmic rays in the high-z Universe 6.2.1 Population III star formation Though CR effects will be examined for a range of star formation rates, the typical Pop III star formation rate is taken to be that found in Bromm and Loeb (2006) at z 15, which is approximately Ψ∗ 2×10 −2 M⊙ yr −1 Mpc −3 in a comoving volume. This rate was derived using the extended Press-Schechter 145
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Copyright by Athena Ranice Stacy 20
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New Insights into Primordial Star F
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Acknowledgments I first want to giv
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angular momentum to rotate at nearl
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2.4.2.5 Thermodynamics of accretion
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Chapter 7. Outlook 177 Bibliography
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List of Figures 2.1 Density project
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1.1 Motivation Chapter 1 Introducti
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leaving behind a neutron star or bl
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describe studies that have attempte
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CRs may therefore provide a pathway
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expected that some of the gas withi
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our numerical methodology in Chapte
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scattering is the major source of o
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portance. The reaction rates for an
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The size of the refinement levels h
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the sink a temperature and pressure
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same time, such as the fragments in
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Figure 2.2: Physical state of the c
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2.4.2 Protostellar accretion 2.4.2.
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Despite some amount of rotational s
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Figure 2.6: Velocity field of the c
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Figure 2.7: Disk mass vs. time sinc
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Figure 2.8: Angular momentum struct
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sink tform [yr] Mfinal [M⊙] rinit
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sink (∼ 4 M⊙) that merged with
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actual final mass of the star is li
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the average temperature of all part
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disk accretion onto a primordial pr
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Figure 2.12: Impact of feedback on
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Figure 2.13: Comparison of specific
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the central core, eventually compri
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extending our simulation to later t
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that such a non-primordial origin t
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impose an upper mass limit to Pop I
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we use a ray-tracing scheme to foll
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3.2.3 Sink Particle Method When an
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convenient way to directly measure
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where the sum now extends from the
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Figure 3.1: Evolution of various pr
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ased on the prescription of Omukai
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ate found at late times in our ‘n
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Figure 3.2: Evolution of various di
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Figure 3.3: Evolution of disk mass
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Figure 3.4: Temperature versus numb
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disk mass is able to level off in d
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Figure 3.6: Projected density and t
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onto the disk. The numerical experi
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that the I-front expands as an hour
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Figure 3.8: Evolution of various H
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˙ M∗ = 3πΣν, (3.18) where ν
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the discrete nature of the gas part
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stars. We also note that, despite t
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Page 111 and 112: It is apparent that the angular mom
Page 113 and 114: h = 0.7. To accelerate structure fo
Page 115 and 116: acc = Lres 50 AU, where: Lres 0.5
Page 117 and 118: years, up to several dynamical time
Page 119 and 120: sp = GM∗ c2 , (4.1) s where cs is
Page 121 and 122: Figure 4.2: Left: Angular momentum
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Page 125 and 126: steadily grows for the rest of the
Page 127 and 128: Fig. 4.5 shows these timescales for
Page 129 and 130: sink B. Once on the MS, the sink A
Page 131 and 132: Figure 4.6: Evolution of stellar ro
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Page 149 and 150: Figure 5.1: Top panels: Effective v
Page 151 and 152: Figure 5.2: Effect of relative stre
Page 153 and 154: Using cs ∝ T 1/2 results in ρ
Page 155 and 156: Figure 5.4: Evolution of gas proper
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Page 173 and 174: β = 1 − −2 1/2 ɛ + 1 mHc2
Page 175 and 176: and where ΓCR(D) = 5 × 10 −29 e
Page 177 and 178: Figure 6.2: Thermal evolution of pr
Page 179 and 180: Figure 6.3: Minimum temperature rea
Page 181 and 182: Figure 6.4: Thermal evolution of pr
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Page 187 and 188: 6.3.5 Fragmentation scale As is evi
Page 189 and 190: 2006). This decrease in fragmentati
Page 191 and 192: Chapter 7 Outlook Our understanding
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Page 197 and 198: Barkana, R. and Loeb, A. (2001). In
Page 199 and 200: Bromm, V., Kudritzki, R. P., and Lo
Page 201 and 202: Daigne, F., Olive, K. A., Silk, J.,
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Machida, M. N., Omukai, K., Matsumo
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Navarro, J. F. and White, S. D. M.
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Rollinde, E., Vangioni, E., and Oli
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Simon, M., Ghez, A. M., Leinert, C.
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Tanvir, N. R., Fox, D. B., Levan, A
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Rays, volume 576 of Lecture Notes i
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Zatsepin, G. T. and Kuz’min, V. A
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