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minihalo to cool to lower temperatures. For CR ionization rates above a critical value of ∼ 10 −19 s −1 , we find that the mass scale of metal-free stars might be reduced to ∼ 10 M⊙, corresponding to what has been termed ‘Pop II.5’. If CRs did indeed facilitate Pop II.5 star formation in z = 20 minihaloes, the subsequent production of CRs would likely have been significantly reduced. Only one Pop III star is expected to form per minihalo, and it is plau- sible that with CRs present only one Pop II.5 star would be able to form in a given minihalo, as feedback effects of this first star may dissipate so much of the minihalo’s remaining gas that no other stars can form in the same halo. A single Pop II.5 star per minihalo would imply no increase in the global number density of stars relative to the case of one Pop III star per minihalo. If Pop II.5 stars ever formed they would all die as lower explosion-energy CCSN. These stars would have lifetimes around 10 7 yr, not significantly longer than for the ∼ 100 M⊙ stars expected to form in minihaloes when CRs are unimportant, so Pop II.5 stars would still generate CRs almost instantaneously. However, their contribution to the CR energy density would be an order of magnitude less than that from a PISN due to the reduced explosion energy. This reduction in overall CR flux might be sufficient to prevent further CR-induced Pop II.5 star formation in minihaloes. Stars formed slightly later could then again be classical Pop III stars which would quickly restore a high level of CR energy density. The overall impact of CRs in the early Universe is thus difficult to determine, owing to the intricate feedback between star formation and CR production. We conclude, however, that CRs could significantly influence primordial star formation, and it will be important to further explore their role with more sophisticated numerical simulations. 176
Chapter 7 Outlook Our understanding of the early universe and the first stars has grown rapidly over just the last several years. The picture of lone, highly massive Pop III stars has now expanded to include disk formation with possible co- formation of low-mass companion Pop III stars. A build-up of a cosmic ray background may also lead to population of relatively small primordial stars. Though results of various simulations still point to a top-heavy IMF, the range of possible Pop III masses is probably much larger than initially thought. Over- all, as computational power and numerical simulations continue to improve, we are finding that the physical processes necessary to describe Pop III star formation are becoming increasingly similar those describing star formation in our own Milky Way. Disk formation, fragmentation, and stellar rotation are important to consider when studying star formation not only at z = 0, but at all redshifts. In Chapter 2, we investigated the formation of metal-free, Pop III, stars within a minihalo at z 20 with a numerical simulation starting from cosmological initial conditions. Beginning from z = 100, we followed the collapsing gas in the center of the minihalo up to number densities of 10 12 cm −3 . This allowed us to study the accretion onto the initial protostellar hydrostatic core, which we represent as a growing sink particle, in improved physical detail. A disk-like configuration assembles around the first protostar, and eventually a 177
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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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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 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: 2006). This decrease in fragmentati
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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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