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are necessary to understand how the relative streaming affects the non-linear regime and alters the processes involved in the collapse of minihalo gas. To this end, we perform a set of cosmological simulations which include these streaming motions. After the completion of this work, we became aware of an analogous paper by Maio et al. (2011). Similar to Maio et al. (2011), we find a delay of gas collapse in early low-mass M ∼ 10 5 − 10 6 M⊙ minihalos, but conclude that for typical streaming velocities this delay will be negligible by z ∼ 10. Our work is complementary to that of Maio et al. (2011) in that, while they are able to find a 1−20% overall suppression of the first objects, our factor of ∼ 10 greater mass resolution allows us to see the subsequent collapse of the gas to high densities, revealing that even with relative streaming motions the thermal evolution of primordial gas and subsequent Pop III star formation will be very similar to no-streaming cases following the initial collapse. 5.2 Numerical Methodology We carry out our investigation using GADGET, a widely-tested three dimensional smoothed particle hydrodynamics (SPH) code (Springel et al. 2001; Springel and Hernquist 2002). Simulations are performed in a periodic box with size of 100 h −1 kpc (comoving) and initialized at zi = 100 with both DM and SPH gas particles. This is done in accordance with a ΛCDM cosmology with ΩΛ = 0.7, Ωm = 0.3, ΩB = 0.04, and h = 0.7. We adopt σ8 = 0.9 for the fiducial normalization of the power spectrum, and also examine the case of σ8 = 1.4 in which structure formation is accelerated and the first minihalo collapses earlier. Each simulation box contains 128 3 DM particles and an equal number of SPH particles. The gas particles each have a mass mSPH = 8 M⊙, so that the mass resolution is, Mres 1.5NneighmSPH < ∼ 400 132
M⊙, where Nneigh 32 is the typical number of particles in the SPH smoothing kernel (e.g. Bate and Burkert 1997). This mass resolution allows us to follow the gas evolution to a maximum number density of nmax = 10 4 cm −3 . The chemistry, heating and cooling of the primordial gas is treated in a fashion very similar to previous studies (e.g. Bromm and Loeb 2004; Yoshida et al. 2006). We follow the abundance evolution of H, H + , H − , H2, H + 2 , He, He + , He ++ , e − , and the deuterium species D, D + , D − , HD, and HD + . We use the same chemical network as used in Greif et al. (2010) and include the same cooling terms. We first perform both the ‘standard collapse’ (σ8 = 0.9) and ‘early collapse’ (σ8 = 1.4) initializations with no streaming velocity added. For each of these we also perform ‘moderate’ and ‘fast’ streaming cases in which we include an initial streaming velocity vs,i of 3 km s −1 and 10 km s −1 , respectively. The ‘moderate’ streaming case represents the predicted root mean square velocity (Tseliakhovich and Hirata 2010), given that peculiar velocities have decreased as (1 + z) since recombination and thus have declined by a factor of 10 at the point our simulations are initialized. Our vs,i values therefore correspond to velocities of 30 km s −1 and 100 km s −1 at recombination, similar to the velocities chosen by Maio et al. (2011), 30 and 60 km s −1 . 5.3 Results 5.3.1 Delay of Gas Collapse The main effect of the relative streaming cases is to delay the collapse of the baryons into the DM halos. In the standard case, the collapse redshifts are zcol = 14.4, 12.2, and 6.6 for vs,i = 0, 3, and 10 km s −1 (0, 30, and 100 km s −1 at recombination). The streaming cases correspond to delays in collapse 133
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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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Page 99 and 100: Figure 3.8: Evolution of various H
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Page 105 and 106: stars. We also note that, despite t
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Page 109 and 110: the deaths of massive stars (see Wo
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
Page 123 and 124: through the disk, and the disk is e
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
Page 133 and 134: sink M∗(5000 yr) [M⊙] M∗(10 5
Page 135 and 136: present here, particularly if the s
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Page 139 and 140: Here, Σ is the disk surface densit
Page 141 and 142: which has already detected GRBs at
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Page 145: are initialized at high z with smal
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 159 and 160: abundances of HD, which also acts a
Page 161 and 162: first stars had very short lifetime
Page 163 and 164: of star-forming material. Here we t
Page 165 and 166: star formation, they found CR energ
Page 167 and 168: The GZK cutoff applies only to thos
Page 169 and 170: and virialization redshift zvir = 2
Page 171 and 172: and undergoes free-fall collapse. I
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
Page 183 and 184: The CR energy density can now be es
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 and 192: Chapter 7 Outlook Our understanding
Page 193 and 194: that Pop III stars likely experienc
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Barkana, R. and Loeb, A. (2001). In
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Bromm, V., Kudritzki, R. P., and Lo
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Daigne, F., Olive, K. A., Silk, J.,
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Gao, L., White, S. D. M., Jenkins,
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Heger, A., Woosley, S. E., and Spru
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Kratter, K. M. and Murray-Clay, R.
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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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