pTPC: advanced gaseous tracking device
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Figure 2. Obtained 3D proton tracks (closed circles) and their Bragg curves shown in
the cathode=O plane. A typical electron track is also shown by open circles.
3. WIMPwind detection with pTPC
Taking into account the tracking capability of the prototype pTPC and
we calculated the demeasured neutron flux at Kamioka Obser~atory'~,
tection sensitivities of WIMPwind8. We assume that the track length
and dElcLx threshold of a p T P C as a WIMPwind detector are 3 inn1 and
10 keV/cm. respectively. F'roni the calculated energy deposition of the F
ion and the scaled track length of the measured va1uel6, we consequently
knew that 25keV F ion has a range of roughly 3mm in 20Torr of CF4.
We also knew that 25 keV Xe ion has a range of roughly 3 mm in 5 Torr of
Xe. The left and right panels of Fig.3 show the SD and SI 3u detection
sensitivities, respectively.
A prototype pTPC as a WIMP detector with a detection volume
30 x 30 x 30cm3 is now being manufactured. Since the fundamental
manufacturing technology is already established, a large volume detector
( w 1ni3) for the underground measurement will soon be available.
4. Conclusion
We found that pTPC filled with CF4 gas is a promising device for the
WIMPwind detection via SD interactions. With even a 0.3 m3. year of
exposure at Kamioka Observatory, it is expected that the best sensitivity
of the current experiments can be achieved. Moreover, it is expected that
the sensitivities of pTPC as WIMP detector can explore the MSSM region
for SD and SI interactions with a sufficient exposure ( 300ni3 . year).
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SD 30 delecllon sensttilies
SI 30 detectionsenstiitles
11
,
4
year
iMSSM
109
10'
102
MWMP [ G W
lo3
Figure 3. Estimated SD (left) and SI (right) 30 detection sensitivities at Kamioka
Observatory for three exposures shown by thick solid lines. Limits from the UKDMC
experiment'? are shown by a thin dotted line (left). Limits from other experiment^'^^^^
are shown by thin dotted lines; and DAMA's allowed region2' is shown by a closed
contour (right). Theory regions predicted by minimal supersymmetric extensions of the
) ~ ~also shown in a thin dotted lines (left and right).
standard model ( I v ~ S S Mare
References
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!
WHAT IS THE REAL ORIGIN OF PRESOLARNOVA GRAINS?*
MARIKO TERASAWA
Centerfor Nuclear Study, Universiw of Tokyo, Hirosawa. Wako. Saitama 351 0198;
mariko@cns.s.utokyo.ac.jp
NOBUYUIU IWAMOTO
Department ojrlstronomy, University of Tokyo. Hongo, Bunkyoku, Tokyo 1130033
We investigate important reactions and reaction paths in order to reproduce the isotopic
ratios of characteristic elements, C, N, and Si, in presolar Sic grains from novae. We fmd
that the Nisotopic ratio strongly depends on the temperature profile in a nova explosion.
By using this temperature dependence, we obtain a favorable temperature profile during a
nova outburst. Moreover, the calculated 3oSi/28Siratio is high compared with the
observational data of presolar nova grains. We also fmd that this overproduction of 30Si
can be avoided if the reaction rate of '@P(~,y)~'s,which is experimentally still unknown,
could increase by a factor of a few tens around the temperature of 3 X 10* K.
1.
Introduction
Grains with specific isotopic ratios different from the solar are called 'presolar
grains'. It has been considered that these presolar grains contain a lot of
information about the site of the nucleosynthesis before the formation of the solar
system ','. Understanding the origin of these grains gives us an insight of not only
the synthesized site but also the chemical evolution of the Galaxy '.
The presolar grains are classified into some groups based on their isotopic ratios.
In this investigation, we specifically focus on the socalled 'nova grains', which
were discovered and reported by Arnari et al. '. Although they inferred a
production site as novae from the fact that the isotopic ratios are "C/13C of 49,
14
N/15N of 520, high '6Al/'7Al, closetosolar 29Si/28Si,and a little excess in
30
Si/"Si, the real origin of nova grains is still questionable. This is because ratios
calculated by hydrodynamical models are largely different from analyzed data,
and the ratios could not be realized by adopting any nova models '. Their results
showed that extra mixing of ejected material with closetosolar matter is needed
more than 95% after nucleosynthesis ended. However, it is also likely that 'nova
grains' come &om only nova ejecta without such a high mixing rate.
* This work is supported by the fellowship of the Japan Society for Promotion of
Science (JSPS).
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542
Therefore, we reexplore reaction paths during nova nucleosynthesis in detail.
As a result, we find the ideal profile of temperature in order to reproduce the
grain data of 12C/13C,I4N/l5N,and 29Si/28Si.It is also showed that the 3oSi/28Si
ratio can be explained, if the reaction rate of an experimentally unknown reaction,
30P(p,y)3'S,changes by a factor of a few tens. Note that we exclude 26Al/27Alratio
from consideration, since the data of 26A1/27A1
have a large uncertainty '.
2.
Calculations and Parameters for Nova Explosions
We adopt the onezone model for nova explosions '. When the temperature at the
bottom of the envelope (Tb)is given, the temperature and density structures in the
envelope are determined by a WD mass ( M w ) and an envelope mass (Men,). We
assume an outburst on an ONe white dwarf (WD) because of the overabundance
of 30Si isotope in nova grains. We change the values of M m between 1.15M,
and 1.35 M,, and M,, between 10j M, and 1O3.0 M,. These parameter ranges
are necessary to eject the envelope '. We have to mention that this onezone
model may not be appropriate near the last phase of the nova outburst, since the
envelope is implicitly assumed to be fully convective. Therefore we assume
temperature and density profiles to decay exponentially during the late phase.
It has been generally assumed that the dredgedup matter from the ONe WD is
mixed with the accreted matter from a companion star. The mixing fraction, XwD,
presents the portion of the dredgedup matter in the mixedmatter. A larger value
of XwD means that novae contain a larger amount of heavy elements. We also
treat the mixing fraction as a parameter, whose ranges are between 0.1 and 0.8.
3.
Results
Nucleosynthesis separately occurs in lower and upper regions of the bottleneck
nuclei with the mass number of A = 19, which have a small proton separation
energy. Each region is related with CNO elements and heavier elements than Ne.
So, we can discuss conditions necessary to reproduce the data of nova grains in
each region, separately. In this paper, we first explain how the temperature profile
is restricted from the temperature dependence of C and Nisotopic ratios and
secondly we show results from Siisotopic ratios (for details see Terasawa and
Iwamoto '). We note that these necessary conditions change slightly with values
of XWD,since different values of X ~ mean
D different initial abundances. Here, we
will show an ideal condition in the case of XwD= 0.8.
In nova explosions, Tb becomes high enough to exceed 2 X 10' K, so that HotCNO cycle (HCNO) occurs. In the HCNO cycle, 15N(p,ct)12Cis the key reaction
to determine C and N ratios because of its high reaction rate. The strong
dependence of the reaction rate on the temperature changes rapidly the "N
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abundance with temperature. Therefore, "N is transformed to 12C under the
condition with relatively high temperature as soon as 1 5 0 decays into "N.On the
other hand, when the temperature is relatively low, "N remains and thus the
14
N/"N ratio becomes low. This dependence of the reaction on the temperature
can give a strong constraint on the profile of temperature evolution after "N is
made by the pdecay of "0, that is, at the late phase of explosions. In order to
leave an appropriate amount of "N,it is needed to synthesize "0 abundantly at
the early phase before the pdecay of "0. For this, it is important that the peak
temperature exceeds 2.8 X 10' K. As we described before, since "0 should
decay in the condition with relatively low temperature below 2 X lo8 K, the
decreasing time from the peak to 2 X 10' K is necessary to be about the plifetime of 1 5 0 , 120 sec. Moreover, it is necessary that the temperature is kept in
a range from 2 X 10' K to 10' K in order to preserve the nuclear flow fiom "N
to "C. The duration is favorable to be several thousands of seconds.
As for Siisotopic ratios, Amari et al. also reported that the 29Si/28Siand
30 . 28 .
Si/ Si ratios are almost the same as the solar. However, both analytical and
hydrodynamical studies 4,' have showed a large enhancement of 30Sirelative to
28
Si by about a factor of 10. In ONe nova models with high peak temperatures,
the reactions which change the abundances of Si isotopes are known to be the
following eight reactions, 2@Ne(p, y)'lNa, 23Na(p,y)24Mg, 23Mg(p,yy4Al,
28Si(p,Y)'~P, 29Si(p,Y)~OP,29P(p,Y)~OS,30P(p,yY1S, and 31P(p,a)28Si'. Since only
the 3oSi/2*Siratio is large compared with grain data, it is favorable that the 30Si
abundance reduces as the 29Siabundance remains unchanged.
Among above eight reactions, the reaction of 30P(p,y)31Shas large and direct
effects on only 30Si abundance This is because the flow from 28Si to heavier
elements does not go through 29Si, and 30Si are made by P'decay of 30P at the
late phase. If the reaction rate of 30P(p,y)3'Sbecomes larger by a few tens, the
flow to heavier elements becomes faster. The remaining abundance of 30P
decreases by a few tens. As a result, 30Siabundance reduces and the Siisotopic
ratios in the nova grains can be reproduced. Moreover, since there has been no
reliable reaction rate for 30P(p,y)31Sdue to the lack of experimental knowledge,
the uncertainty is still a factor of 100 up and down '. Accordingly, it is quite
possible that the reaction rate could be a factor of a few tens higher than the
current rate.
Thus, when its reaction rate becomes higher by a few tens at the temperature
around 3 X 10' K, a good fit is obtained with Siisotopic ratios measured in nova
grain candidates. For example, we can see a good agreement, 12C/13C= 7.74 and
14
N/"N = 5.56, in the case of an explosion with MwD= 1.3M,, M,, = 2.5 X
M,, and XwD= 0.8. The reaction rate of 30P(p,y)31Sis multiplied by 20.
At the end, we describe shortly the dependence of temperature profiles on the
value of XwD.We find that it is necessary for the temperature to reach 3.5 X lo8
K in order to make sufficient 29Si in the case of XWD = 0.4. This constraint is
more rigorous than that from the Nisotopic ratio. Thus, synthesis of Si isotopes
'
'.
544
imposes a constraint on the peak temperature. When a larger value of XwD is
adopted, the initial abundance of "Si becomes higher. Then, the peak
temperature is allowed to be low. Actually, the peak temperature of 2.8 X 10sK
suffices in the case of X m = 0.8.

4.
Discussions
We have to mention that the most remarkable difference between
hydrodynamical models and our onezone models is made at the late phase of
outbursts because of different treatments for convection. However we recognize
that abundances of 12C, "C and ''N strongly depend on the temperature profile at
the late phase. Previous hydrodynamical simulations showed the lower values of
C and Nisotopic ratios than nova grains '. Based on their results and our results
in which temperature remains relatively high even at the late phase, the inner and
narrow regions in the envelope may be preferable to form the nova grains.
Therefore, nova grains may be seldom found though nova fiequency is relatively
high in our Galaxy (30+ 10 yrI).
Acknowledgments
One of the authors (MT) would like to thank Prof. S. Kubono for useful
discussions.
References
1. E. Anders and E. Zinner, Meteoritics, 28,490 (1993)
2. E. Zinner, Meteoritics and Planetary Science, 33,549 (1998)
3. T. J. Bernatowicz and R. Cowsik, American Institute o j Physics Conference
Series, 402,45 1 (1997)
4. S. Amari, X. Gao, L. R. Nittler, E. Zinner, J. Jose, M. Hernanz and R. S. Lewis,
ApJ, 551,1065 (2001)
5. S. Wanajo, M. Hashimoto and K. Nomoto, ApJ, 523,409 (1999)
6. M. Terasawa and N. Iwamoto, submitted to ApJL (2004)
7. C. Iliadis, A. Champagne, J. Jose, S . Starrfield, and P. Tupper, ApJS, 142, 105
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8. J. Jose, A. COC,and M. Hernanz, ApJ, 560,897 (2001)
AMD+GCM STUDY OF STRUCTURE OF CARBON ISOTOPES*
G. THIAMOVA~,N. ITAGAKI, T. OTSUKA
Department ofphysics, Universiv of Tokyo, Hongo, Tokyo 1130033, Japan
K. IKEDA
The Institute of Physical and Chemical Research (RIKEN), Wako,
Saitama, 3510198, Japan
The ground state properties of the carbon isotopes are inveshgated usmg the extended
version of the Antisymmetrized Molecular Dynamcs (AMD) Multl Slater Deternunant
method We can reproduce reasonably well many expenmental data for 12C22CIn this
contribuhon we present a systematic calculation of bmdmg energles, energies of the 2'+
states and B(E2) transition strengths
1. Introduction
The AMD method is very suitable for the description of light systems where
both shellmodel and cluster structures can appear because it is free from any
model assumption concerning the wave functions.
The extended version of the AMD method adopted in this work corresponds
to the combination of AMD and the Generator Coordinates Method (GCM)
[l]. The initial GCM basis functions are prepared in such a way that they
correspond to several properly chosen r.m.s. radii constraints, close to the
experimental values.
The mixing amplitudes of these Slater determinants are determined by diagonalization of the Harniltonian matrix. In this way suitable basis for the GCM
calculations can be obtained. The theoretical details of the method are explained
in [2].
2. Results
The Hamiltonian and the effective nucleonnucleon interaction used is the same
as in [3]. The calculations are performed with 45 and 60 basis functions for
*
This work is supported by Grantin Aid for Scientific Research (13740145)
and by The Japanese Society for Promotion of Science under the contract No
On leave of absence from the Nuclear Physics Institute, Czech Academy of
Sciences, PragueRez, Czech Republic.
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546
eveneven and evenodd isotopes, respectively. The details concerning the basis
hnctions can be found in [2].
The binding energies are presented in Fig. 1. In general, good agreement is
obtained in all the studied region The binding energy of 12C is smaller than the
experimental value. It is partially due to the Majorana parameter M=0.6, fitted
to the binding energy of l 6 0 and known to produce underbinding of I2C. On the
other hand, the spinorbit term seems to be too strong and thus the 3alpha
component in the ground state wave function is too small. This is also reflected
in the smaller B(E2) transition strength (see below).
To describe a halo nucleus I5C is a real challenge for the AMD methods.
Here we do not reproduce the ground state spin 1/2+ . This is mainly due to the
simple interaction with no tensor term and strong spinorbit term. However, in
[2] we have adopted a better description of the sorbit for the odd neutron and
the excitation energy of the 112' decreased considerabely.
The systematics of the excitation energies of the 2'1 states clearly supports
the idea about N=16 magic number, reflected by large 2fl energy of 22C. The
(dj/2)6subshell closure predicted by our calculation but not seen experimentally
is again due to the stronger spinorbit term, which pushes the dj/2 orbit down in
energy. A comparison is made with an AMD calculation [4] with weaker spinorbit term and modified Volkov interaction W 1 .
The B(E2 ) transition strengths (Fig.3) are compared with the experimental
data and the shellmodel values [5] obtained with effective charges. Smaller
B(E2) value for I2C reflects most probably smaller 3alpha component in the
ground state wave function due to stronger spinorbit term. In I6,l8C protons
construct almost closed shellmodel configuration so the B(E2) value is very
small. Proton contribution is recovered again in *OC. The very small B(E2) value
for I6C has been measured recently [6] and is successfully reproduced by our
model.
3. Summary
We have performed a systematic AMD+GCM calculation of structure of carbon
isotopes 12C**C.We can reproduce fairly well a lot of experimental data. Here
we present the systematic calculation of binding energies, 2cI energies and
B(E2) strengths. Even though the effective interaction is simple and there are
indications that the spinorbit term is too strong it should not change the
qualitative results of this analysis. From the systematics of 2+1 energies a clear
support for the N=16 magic number is given. B(E2) value of 12C is smaller due
to strongerspin orbit term. Very small B(E2) value for I6C is successhlly
reproduced by our model.