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small multiple system forms, dominated by a binary with masses ∼ 40 M⊙ and ∼ 10 M⊙. If Pop III stars were to form typically in binaries or small multiples, this would have crucial consequences for the observational signature of the first stars, such as the nucleosynthetic pattern of the SNe they produce, and the gravitational-wave emission from possible Pop III black-hole binaries. In Chapter 3, we examined the growth of metal-free, Pop III stars under radiative feedback. Using the sink particle method within a cosmological simulation, we modeled the effect of LW radiation emitted by the protostar, and employed a ray-tracing scheme to follow the growth of the surrounding H ii region. We again find that a disk assembles around the first protostar, and that radiative feedback will not prevent further fragmentation of the disk to form multiple Pop III stars. The resuling heating and ionization of the gas leads to a shutoff of accretion once the main sink has grown to 15 M⊙ while the second sink has grown to 5 M⊙. There is little further accretion or fragmentation afterwards, indicating the most likely outcome is a massive Pop III binary. If Pop III stars were typically unable to grow to more than a few tens of solar masses, this would have important consequences for the occurence of pair-instability supernovae in the early universe as well as the Pop III chemical signature in the oldest stars observable today. In Chapter 4 we employed a cosmological simulation to estimate the rotation speed of Pop III stars within a minihalo. Again using sink particles, we measured the velocities and angular momenta of all particles that fall onto these protostellar regions. This allowed us to record the angular momentum of the sinks and estimate the rotational velocity of the Pop III stars expected to form within them. We find that there is sufficient angular momentum to yield rapidly rotating stars (> ∼ 1000 km s −1 , or near break-up speeds). This indicates 178
that Pop III stars likely experienced strong rotational mixing, impacting their structure and nucleosynthetic yields. A subset of them was also likely to result in hypernova explosions, and possibly GRBs. In Chapter 5, we evaluated how the recently discovered supersonic relative velocity between the baryons and dark matter will affect the thermal and density evolution of the first gas clouds at z < ∼ 50. We find that the typical streaming velocities will have little effect on the gas evolution. Once the collapse begins, the subsequent evolution of the gas will be nearly indistinguishable from the case of no streaming, and star formation will still proceed in the same way, with no change in the characteristic Pop III stellar masses. Reionization is expected to be dominated by luminous sources that form within DM halos of masses > ∼ 10 8 M⊙, for which the effect of streaming should be negligible. In Chapter 6, we explored the implications of a possible CR back- ground generated during the first SNe explosions that end the brief lives of massive Pop III stars. CRs are expected to propagate away from SNe much faster than the metals produced at the same site. We show that a CR back- ground could have significantly influenced the cooling and collapse of primordial gas clouds in minihalos around redshifts of z ∼ 15 − 20, provided the CR flux was sufficient. The presence of CRs could indirectly enhance the molecular cooling in these regions, and we estimate that the resulting lower temperatures in these minihalos would yield a characteristic stellar mass as low as ∼ 10 M⊙. Future work will aim to accomodate this developing picture of Pop III star formation in simulations of first galaxy formation. The metallicity and radiation produced by the first stars, and thus the environment in which future generations of stars formed within the first galaxies, depends crucially on the 179
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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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current computational limits. If ra
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the deaths of massive stars (see Wo
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It is apparent that the angular mom
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h = 0.7. To accelerate structure fo
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acc = Lres 50 AU, where: Lres 0.5
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years, up to several dynamical time
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sp = GM∗ c2 , (4.1) s where cs is
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Figure 4.2: Left: Angular momentum
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through the disk, and the disk is e
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steadily grows for the rest of the
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Fig. 4.5 shows these timescales for
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sink B. Once on the MS, the sink A
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Figure 4.6: Evolution of stellar ro
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sink M∗(5000 yr) [M⊙] M∗(10 5
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present here, particularly if the s
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the stars grow to larger masses (se
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Here, Σ is the disk surface densit
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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
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Page 155 and 156: Figure 5.4: Evolution of gas proper
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Page 163 and 164: of star-forming material. Here we t
Page 165 and 166: star formation, they found CR energ
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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 185 and 186: 2007). When considering realistic c
Page 187 and 188: 6.3.5 Fragmentation scale As is evi
Page 189 and 190: 2006). This decrease in fragmentati
Page 191: 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.,
Page 203 and 204: Gao, L., White, S. D. M., Jenkins,
Page 205 and 206: Heger, A., Woosley, S. E., and Spru
Page 207 and 208: Kratter, K. M. and Murray-Clay, R.
Page 209 and 210: Machida, M. N., Omukai, K., Matsumo
Page 211 and 212: Navarro, J. F. and White, S. D. M.
Page 213 and 214: Rollinde, E., Vangioni, E., and Oli
Page 215 and 216: Simon, M., Ghez, A. M., Leinert, C.
Page 217 and 218: Tanvir, N. R., Fox, D. B., Levan, A
Page 219 and 220: Rays, volume 576 of Lecture Notes i
Page 221 and 222: Zatsepin, G. T. and Kuz’min, V. A
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