Proton fraction in neutron stars

Abstract

The proton fraction in ${\sl \beta}$-stable neutron stars is investigated within the framework of the Skyrme-Hartree-Fock theory using the extended Skyrme effective interaction for the first time. The calculated results show that the proton fraction disappears at high density, which implies that the pure neutron matter may exist in the interior of neutron stars. The incompressibility of the nuclear equation of state is shown to be more important to determine the proton fraction. Meanwhile, it is indicated that the addition of muons in neutron stars will change the proton fraction. It is also found that the higher-order terms of the nuclear symmetry energy have obvious effects on the proton fraction and the parabolic law of the nuclear symmetry energy is not enough to determine the proton fraction. Comment: 4 Pgaes in REVTex, 2 Figures, 1 Table

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arXiv:nucl-th/0011017v1 6 Nov 2000
Proton Fraction in Neutron Stars
Feng-Shou Zhang1,2,4, Lie-Wen Chen1,2,3
1Center of Theoretical Nuclear Physics, National Laboratory of Heavy Ion
Accelerator, Lanzhou 730000, China
2Institute of Modern Physics, Chinese Academy of Sciences, Lanzhou 730000, China
3Department of Applied Physics, Shanghai Jiao Tong University, Shanghai 200030, China
4CCAST (World Laboratory), P.O. Box 8730, Beijing 100080, China
The proton fraction in β-stable neutron stars is investigated within the framework
of the Skyrme-Hartree-Fock theory using the extended Skyrme effective interaction
for the first time. The calculated results show that the proton fraction disappears at
high density, which implies that the pure neutron matter may exist in the interior of
neutron stars. The incompressibility of the nuclear equation of state is shown to be
more important to determine the proton fraction. Meanwhile, it is indicated that the
addition of muons in neutron stars will change the proton fraction. It is also found
that the higher-order terms of the nuclear symmetry energy have obvious effects on the
proton fraction and the parabolic law of the nuclear symmetry energy is not enough
to determine the proton fraction.
PACS: 97.60.Jd, 26.60.+c
The neutron star serves as a cosmic laboratory to study the nuclear equation of state (EOS) at high
density. It has been known that the mass, radius, and density profile of neutron star are determined by
the EOS of asymmetric nuclear matter, and the chemical composition, particularly the proton fraction
is determined by the nuclear symmetry energy1. Many theoretical methods, such as the Brueckner
approach2, variational many body (VMB) theory3, the relativistic mean field (RMF) theory4, and
non-relativistic many body theory5, have been used to study the proton fraction in neutron stars
since the proton fraction is an important physical quantity to determine the evolution of neutron
stars, such as the cooling rate, neutrino flux, and so on6−9. However, little is known about the
properties of the nuclear symmetry energy at high density10with the proton fraction in neutron stars
remaining undetermined.
In our previous work11,12, a density, temperature and isospin dependent nuclear EOS has been
deduced within the framework of the Skyrme-Hartree-Fock theory using the so-called extended Skyrme
effective interaction13. For asymmetric nuclear matter with density ρ and neutron excess δ = (ρn−
ρp)/ρ , the energy per nucleon at temperature T can be written as
ε(ρ,T,δ) =1
2T
?C3/2(µτα)
C1/2(µτα)(1 + δ)5/3+C3/2(µτα)
C1/2(µτα)(1 − δ)5/3
?
+
1
4
1
4
1
4
1
4
?
?
?
?
a(1)
τα(1 + δ) + a(1)
−τα(1 − δ)
?
?
?
?
ρ +
a(2)
τα(1 + δ) + a(2)
−τα(1 − δ)ργ+1+
a(3)
τα(1 + δ) + a(3)
−τα(1 − δ)ρ5/3+
a(4)
τα(1 + δ) + a(4)
−τα(1 − δ)ργ+5/3,(1)
with
a(1)
τα=1
4t0[3 ∓ (2x0+ 1)δ],
a(2)
τα=
1
24t3[3 ∓ (2x3+ 1)δ],
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a(3)
τα=
1
16π2[t1(1 − x1) + 3t2(1 + x2)](1 ± δ)5/3
?2√π
?2√π
λ
?5
?5
C3/2(µτα)
+
1
8π2
?
t1(1 +x1
2) + t2(1 +x2
2)
?
(1 ∓ δ)5/3
λ
C3/2
?µ−τα
?,
a(4)
τα=
1
16π2[t4(1 − x4) + 3t5(1 + x5)](1 ± δ)5/3
?2√π
λ
?2√π
?5
C3/2(µτα)
+
1
8π2
?
t4
?
1 +x4
2
?
+ t5(1 +x5
2)
?
(1 ∓ δ)5/3
λ
?5
C3/2(µ−τα).
where ταand −ταrepresent neutron and proton, respectively. Every equation is split into two equa-
tions, which are for the neutron and proton, respectively. The t0−t5and x0−x5are Skyrme potential
parameters, λ and Cl(µτα) represent the thermal wave length and Fermi-Dirac integral.The Taylor ex-
pansion of Eq. (1) with respect to δ at temperature T = 0 (hereafter,we define ε(ρ,δ) = ε(ρ,δ,T = 0))
is as follows
ε(ρ,δ) = ε(ρ,δ = 0) + C(ρ)δ2+ D(ρ)δ4+ O(δ6), (2)
where C is the so-called coefficient of symmetry energy strength.
The chemical composition of the neutron star is determined by the requirement of charge neutrality
and equilibrium with respect to the weak interaction ( β-stable matter ). From Eq. (1) or (2) we can
calculate the proton fraction, xp= (1 − δ)/2, for β-stable matter as found in the interior of neutron
stars. Equilibrium for the reactions n → p + e−+ νeand p + e−→ n + νerequires that
µn= µp+ µe,(3)
where µi= ∂ε/∂Ni(i = n, p, e and µ) is the chemical potential. Consequently, one obtains
µn− µp= 2∂ε
∂δ.
(4)
For relativistic degenerate electrons,
µe= (m2
e+ k2
Fe)1/2=
?
m2
e+ (3π2ρxe)2/3?1/2
,(5)
with ¯ h = c = 1 and xp= xebecause of charge neutrality. From Eqs. (1), (3)-(5), it is easy to find
numerically the proton fraction as a function of density ρ.
Just above nuclear matter density at which µe exceeds the muon mass mµ, the reactions e−→
µ−+ νe+ νµ, p + µ−→ n + νµand n → p + µ−+ νµmay be energetically allowed such that both
electrons and muons are present in β-stable matter. This alters the β-stability condition to
µn− µp= µe= µµ=
?
m2
µ+ (3π2ρxµ)2/3?1/2
,(6)
with xp= xe+ xµ
The bulk properties of symmetry nuclear matter at saturation density, such as the density ρ0, the
binding energy per nucleon ε0, the incompressibility K, the effective mass m∗/m, and the symmetry
energy are summarized in Table 1 for the three parameter sets SII, SIII, and SkI5. From Table 1, one
can see that the difference for incompressibility between SII and SkI5 is very large, about 85 MeV
at saturation density, while they almost have the same m∗/m=0.58. Meanwhile, the difference for
m∗/m is very large between SIII and SII, about 0.18 at saturation density, while the difference for
incompressibility between them is relatively small, about 14 MeV. In calculations, we note that the
binding energy per nucleon ε and the nuclear symmetry energy C exhibit completely different behaviors
with the increase of the baryon density for different Skyrme parameter sets. These differences come
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from the difference of the nuclear EOSs. Therefore, one may investigate the proton fraction for
different nuclear EOSs.
Under the parabolic approximation (up to the -δ2term in Eq. (2)), the chemical compositions of
neutron stars in both the electrons only and electrons plus muons cases as a function of the baryon
density are shown in Figs. 1 (a)-1 (c) using the Skyrme parameter sets SII, SIII, and SkI5, respectively.
In order to ensure validity of the proton fraction above, we show the squared ratio of sound speed
to light speed, (s/c)2= ∂P/∂E (where E is the total energy density), in Fig. 1 (d). The causality
condition requires (s/c)2≤ 1. One notes that for SII and SkI5, the squared ratio (s/c)2will make
the causality condition fail to be satisfied at baryon density larger than 1.35 fm−3and 1.1 fm−3,
respectively, about 8 times larger than the saturation density for symmetry nuclear matter. For SIII,
up to baryon density 1.5 fm−3, the causality condition can be still satisfied. From Figs. 1 (a)-1 (c),
one can find that the addition of muons increases significantly the proton fraction over a broad range
of density. In case of SIII, the muon fraction is very small and muons only exist at about 0.16 fm−3,
which is due to the fact that µeis too small to exceed the muon mass mµ. The value of the proton
fraction is an important ingredient to model the thermal evolution of a neutron stars. In fact, if the
proton fraction in the interior of star is above a critical value xUrca, the so-called direct Urca processes
n → p + e−+ νe and p + e−→ n + νe can occur14. If they occur, neutrino emission and neutron
star cooling rate are increased by a large factor compared to the standard cooling scenario14. The
critical proton fraction has been estimated to be in the range of xUrca= 11 − 15%. It is indicated
from Fig. 1 that the proton fraction for SII and SIII cannot reach this range. However, the proton
fraction reaches xUrca= 11% at baryon density 0.2 fm−3for SkI5. This means that one can get a
reasonable proton fraction for a softer nuclear EOS, such as K = 255 MeV for case of SkI5. While
for stiffer ones, such as SII (K = 341 MeV) and SIII (K = 355 MeV), one might not get reasonable
proton fraction although they may have a very different effective mass (m∗/m = 0.580 and 0.763 for
SII and SIII, respectively). It seems that the incompressibility is more important to determine the
proton fraction.
One of the interesting features of Fig. 1 is the disappearance of the proton fraction at high density,
namely, ρ = 0.48 and ρ = 0.32 fm−3for SII and SIII, respectively. This phenomenon may be due
primarily to the greater short-range repulsion in isospin singlet nucleon pairs compared to isospin
triplet nucleon pairs15. At high density this short-range repulsion must dominate and pure neutron
matter is favored. At intermediate densities the strong isospin singlet tensor potential and correlation
serve to keep the isospin singlet nucleon pairs, and thus symmetry nuclear matter, more attractive
than pure neutron matter. These features imply that the pure neutron matter may exist in the interior
of neutron stars where the density is high enough. This phenomenon has also been observed in the
calculations of the VMB theory. For parameter set SkI5, the disappearance of the proton fraction is
not found up to 1.5 fm−3.
Many theoretical studies have shown that the empirical parabolic law of the single particle binding
energy is well satisfied within a broad range of density and neutron excess δ, namely, in Eq. (2) the
δ2term is significant, but the higher-order terms are very small. However, it is still unknown how
the higher-order terms affect the proton fraction at high density. Fig. 2 is the same as Fig. 1 except
for including all the higher-order terms. From the comparison between Figs. 1 and 2, one can find
that the addition of the high-order terms increases significantly the fraction of different constituents,
p, e−and µ−, over a broad range of density for parameter sets SII and SIII. It implies that the
higher-order terms of the nuclear symmetry energy have obvious effects on the chemical composition
and the parabolic approximation of the nuclear symmetry energy may result in large deviation.
The authors would like to thank Professor Peng Qiu-He for interesting discussions. This work was
supported by the National Natural Science Foundation of China under Grant Nos. 19875068 and
19847002, the Major State Basic Research Development Program under Contract No. G2000077407,
and the Foundation of the Chinese Academy of Sciences.
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TABLE CAPTIONS
Table I Bulk properties of symmetry nuclear matter at saturation point.
Parameter
ρ0 (fm−3)
ε0 (MeV/nucleon) -15.99 -15.85 -15.85
K (MeV)341.38 355.35 255.78
m∗/m0.580
C (MeV)34.26
SIISIIISkI5
0.1484 0.1453 0.1558
0.763
28.25
0.579
36.75
FIGURE CAPTIONS
Fig. 1 With parabolic approximation, the compositions of neutron stars in both the electrons only
and electrons plus muons cases as a function of the baryon density using the Skyrme parameter
sets SII (a), SIII (b), SkI5 (c). Meanwhile, the squared ratio of sound speed to the light speed
(s/c)2as a function of the baryon density for these three parameter sets is displayed in (d) in
electrons plus muons case. In (b), the muon fraction is very small so that the solid line, dotted
line and dash-dotted line coincide with each other.
Fig. 2 Same as Fig. 1 but including all higher-order terms instead of parabolic approximation for
nuclear symmetry energy.
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0.00.3
Baryon density (fm-3)
0.60.91.21.5
0.0
0.3
0.6
0.9
1.2


(d)
ε=ε(ρ,0)+C(ρ)δ
p (n,p,e-,µ-)
e- (n,p,e-,µ-)
µ- (n,p,e-,µ-)
p (n,p,e-)
2
Fraction
(s/c)
2
FIG.1



0.0
0.1
0.2
0.3
0.4

(c) SkI5


0.00
0.01
0.02
0.03
0.04

(b) SIII



0.00
0.02
0.04
0.06
0.08

(a) SII
(n,p,e-,µ-)
SII
SIII
SkI5