Dust Evolution in Population III Supernova Remnants

Takaya Nozawa , Takashi Kozasa, Asao Habe , Eli Dwek† , Hideyuki Umeda ,
Nozomu Tominaga , Keiichi Maeda‡ and Ken’ichi Nomoto ,§Department of Cosmosciences, Graduate School of Science, Hokkaido University, Sapporo 060-0810, Japan†
Laboratory
for Astronomy and Solar Physics, NASA Goddard Space Flight Center, Greenbelt, MD 20771
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Department of Astronomy, School of Science, University of Tokyo, Bunkyo-ku, Tokyo 113-0033, Japan‡
Max-Planck-Institut für Astrophysik, Karl-Schwarzschild Strasse 1, 85741 Garching, Germany§
Research Center for the Early Universe, School of Science, University of Tokyo, Bunkyo-ku, Tokyo 113-0033,
Japan
Abstract. We present the results of calculations for the evolution of dust within Population III supernova remnants (SNRs),
focusing on the dust formed in the unmixed ejecta of Type II SNe. We show that once dust grains inside the He core encounterthe reverse shock, they are subject to different fates depending on their initial sizes aini . For SNRs expanding into the interstellar medium (ISM) with nH  0  1 cm 3 , grains of aini 0 05 µ m are trapped in the hot gas and are completely destroyed;grains of aini 0.05–0.2 µ m are trapped in the dense shell behind the forward shock with the final sizes of 0.001–0.1 µ m;
grains of aini 0 2 µ m are injected into the ISM without significant destruction. The total mass of surviving dust is 0.01 to
0.8 M and is higher for the lower ISM gas density. We also investigate the elemental abundances of the second-generationstars that form in the dense shell of Population III SNRs, based on the elemental composition of dust piled up in the shell.
The comparison of those results with the observations of hyper-metal-poor (HMP) stars indicates that the transport of dust
segregated from metal-rich gas within a SNR can be responsible for the abundance patterns of Mg and Si in HMP stars.
Keywords: dust, extinction — early universe — shock waves — supernova remnants — supernovae: general
PACS: 95.30.Lz, 95.30.Wi, 97.10.Tk, 97.20.Wt, 97.60.Bw, 98.38.Mz

INTRODUCTION
Dust grains in the early universe play crucial roles in
the formation processes of stars and galaxies through thecooling of gas by their thermal emission ([1, 2]) and theformation of H2 molecules on their surfaces ([3, 4]). In
addition, the extinction of stellar light by dust causes a
serious misleading in interpreting the star formation rate
and the initial mass function of the very early generationof stars from the observations toward high redshifts. The
extinction and thermal emission by dust grains strongly
depend on their chemical compositions and size distributions. Thus, in order to elucidate the evolution of theuniverse from both theoretical and observational studies,
it is indispensable to reveal the properties of dust in theearly epoch of the universe.
Given the cosmic age less than 1 Gyr, the main sourcesof dust are considered to be supernovae (SNe) [5], since
the dust formation in stellar winds from evolved lowmass stars requires too long timescale to supply copious
amounts of dust to the interstellar medium (ISM). Some
theoretical studies ([6, 7, 8, 9]) have investigated thecomposition, size, and mass of dust formed in the ejecta
of primordial Type II SNe (SNe II) and pair-instability
SNe. However, the newly condensed dust grains are subsequently processed by the reverse shock generated fromthe interaction of the SN ejecta with the surrounding

medium ([9, 10]). Therefore, the properties and amountof dust injected from SNe into the ISM significantly differ from those at the time of dust formation.
In this proceedings, we present the results for the evolution of the dust formed in primordial SNe by examining its processing through the collisions with the reverse
shocks and its transport within SN remnants (SNRs). Our
main aim is to reveal how much amount of dust can survive the reverse shock and its dependence on the progenitor mass and the ambient gas density. It is shown that apart of dust grains surviving the destruction are trappedin the dense SN shell, which may have significant impacts on the abundance patterns of the stars formed there.
Thus, we investigate the elemental abundances of thesecond-generation stars formed in the dense shell of Population III SNRs and compare those with the observations of hyper-metal-poor (HMP) stars.

DUST EVOLUTION IN POP III SNRS
The details of calculations of the dust evolution in Population III SNRs are described in [10], where the dynamics and destruction of dust are carefully treated by taking
into account the size distribution and the initial spatial
distribution of each dust species. In [10], the time evolution of gas temperature and density in SNRs is calculated

for the uniform ISM with the gas temperature of 104 K
and hydrogen number density of nH 0 0.1, 1, and 10
cm 3 , based on the hydrodynamic models of Population
III SNe by [11]. For the models of the initial dust within
the He core, the results of the dust formation calculations by [7] are adopted. In this proceedings, we present
the evolution of the dust formed in the unmixed ejecta of
SNe II whose progenitor masses are Mpr 13, 20, 25,
and 30 M and explosion energy is 1051 ergs.
Fig. 1a shows the trajectories of C, Mg2 SiO4 , and Fe
grains in the SNR for Mpr 20 M and nH 0 1 cm 3 ,and Fig 1b represents the time evolution of their sizes. In
Fig. 1a, the trajectories of the forward shock, the reverse
shock, and the surface of the He core are also depictedby the thick solid lines. C, Mg2 SiO4 , and Fe grains inside the He core intrude into the reverse shock at 3650
yr, 6300 yr, and 13000 yr, respectively, and their transport and survive heavily depend on their initial sizes ainiand compositions. Grains with aini 0 01 µ m (dottedlines) are efficiently decelerated by the gas drag and arequickly trapped in the hot gas of 106 K between theforward and reverse shocks. These small grains continue
to be eroded by the thermal sputtering until they are completely destroyed. For grains with aini 0 1 µ m (solidlines), the kinetic and/or thermal sputterings reduce theirsizes, but the grains can survive without being completely destroyed. C and Mg2 SiO4 grains of aini 0 1µ m are finally trapped in the dense SN shell formed at
2 105 yr with the sizes of 0.057 and 0.031 µ m, respectively. In the dense shell, the erosion of dust by thethermal sputtering does not work because the gas temperature drops down rapidly below 105 K ([12]). Fe grainsof aini 0 1 µ m are injected into the ISM because of its
high bulk density, with the size of about half the initialone. Grains with aini 1 µ m (dashed lines) can maintaintheir high velocities and pass through the forward shock
to be injected into the ISM, undergoing the kinetic sputtering in the postshock flow; their sizes decrease only by
0.7%, 6%, and 8% of the initial ones for C, Mg2 SiO4 ,and Fe grains, respectively.
The behavior of the transport and destruction of dust
within a SNR is almost independent of the progenitor
mass considered here; for nH 0
1 cm 3 , grains withaini 0 05 µ m are completely destroyed by sputteringin the hot gas, grains with aini 0.05–0.2 µ m are piled
up in the dense SN shell with the final sizes of 0.001–0.1µ m, and grains with aini 0 2 µ m are ejected into the
ISM without being processed significantly. Accordingly,
the size distribution of the surviving dust in mass isgreatly populated by large sizes, compared with that at
its formation. The total mass fraction of surviving dust
somehow depends on the progenitor mass (Fig. 2), since
the size distribution of each dust species at the time ofdust formation is different from model to model.






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FIGURE 1. (a) Trajectories of C, Mg2 SiO4 , and Fe grains
within the SNR for Mpr  20 M and nH  0  1 cm 3 and (b)ratios of their sizes to the initial ones as a function of time. The
evolution of dust with aini  0.01, 0.1, and 1 µ m is denoted by
the dotted, solid, and dashed lines, respectively. The thick solidlines in (a) indicate the trajectories of the forward and reverse
shocks and the position of the surface of the He core.



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FIGURE 2. Total mass fraction of surviving dust vs. theprogenitor mass. The results for nH  0  0.1, 1, and 10 cm 3 areconnected by the dotted, solid, and dashed lines, respectively.





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On the other hand, the ambient gas density strongly
affects the evolution of dust in a SNR because the higher
ISM density results in the higher density of the shockedgas and causes the efficient erosion and deceleration ofdust due to more frequent collisions with the hot gas.
Thus, with increasing the density in the ISM, the initialsize below which dust is completely destroyed increases(0.01, 0.05, and 0.2 µ m for nH 0 0.1, 1, and 10 cm 3 ,respectively), and the total mass fraction of the surviving
dust is depressed as shown in Fig. 2. The total mass ofsurviving dust is 0.01–0.8 M for nH 0 10 to 0.1 cm 3 .








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ELEMENTAL ABUNDANCES OF THE
SECOND-GENERATION STARS

5 3 ([13,Two HMP stars, HE 0107-5240 with [Fe/H]14]) and HE 1327-2326 with [Fe/H]
5 45 ([15, 16]),
show extreme (>102 times) overabundances of C, N, and
O and large (1–100 times) enhancements of Mg and Si
relative to Fe. Recently discovered HE 0557-4880 with
[Fe/H]
4 75 ([17]) is also a carbon-rich ultra-metalpoor (UMP) star. Although the elemental compositionsof these low-mass stars are expected to reflect the nucleosynthesis in Population III stars, the origin of the peculiar abundance patterns in HMP and UMP stars is still amatter of controversy.
The results for the evolution of dust within a SNR
given in this paper show that the dust grains which survive the passage of the reverse shock but are not injected
into the ISM are accumulated in the dense SN shell in
105 –106 yr. This transport of dust to the SN shell may
enable the formation of stars with solar mass scales there([1, 2]). Thus, we investigate whether the metal abundance patterns of HMP and UMP stars can be explainedby the elemental composition of dust piled up in theshell, assuming that they are the second-generation stars
formed in the shell of Type II SNRs.
We present the abundances of C and O in the denseshell in Fig. 3a and those of Mg and Si in Fig. 3b,
derived from the results of calculations for Mpr 20 M
(filled symbols). The calculated [Fe/H] is 5 20, 5 53,
and 4 92 for nH 0 0.1, 1, and 10 cm 3 , respectively,
which are in good agreement with those for HMP and
UMP stars. In addition, the abundances of Mg and Si for
nH 0 0.1 and 1 cm 3 are in the range of 0 [Mg, Si/Fe]
2 and reproduce the modest overabundances of Mg
and/or Si in HMP stars. Other SN models with Mpr 13,
25, and 30 M also give the similar results for [Fe/H] and[Mg, Si/Fe] ([10]). Therefore, we can conclude that thetransport of dust segregated from metal-rich gas within
a SNR can be an important process in determining theabundance patterns of refractory elements such as Mg,
Si, and Fe in the second-generation stars that form in thedense shell of primordial SNRs. It should be noted thatthe extreme excesses of C and O as observed in HMP
stars can not be reproduced by the models consideredhere. Then, we examine the abundance patterns in theshell by assuming that not only the piled-up grains but
also the gas outside the innermost Fe layer of the SN
is incorporated into the shell. This case can produce theextreme ( 100 times) overabundances of C and O but
leads to unreasonable excesses of Mg and Si as well(see the open symbols in Fig. 3). However, it might bepossible to reproduce the elemental abundances of HMP
stars if the Si-Mg-rich layer is not mixed with the gas in
the shell. This subject will be left for the future work.


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FIGURE 3. Abundances of (a) C and O and (b) Mg and Si in
the shell of the SNR for Mpr  20 M . The filled symbols arethose from the elemental compositions of the grains piled up
in the shell, and the open symbols are those including the gas
outside the innermost Fe layer in addition to the piled-up grains.
The observational data for HMP and UMP stars are taken from
[14, 15, 16, 17, 18] and are denoted by the double symbols.

ACKNOWLEDGMENTS



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This work has been supported in part by a Grant-in-Aid
for Scientific Research from the Japan Society for the
Promotion of Sciences (19740094 and 18104003).

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