neutron stars: probing ultra-dense and hot...
TRANSCRIPT
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Neutron stars: probing ultra-dense andhot matter
Micaela [email protected]
Laboratoire Univers et Théories (LUTH)CNRS / Observatoire de Paris/ Université Paris Diderot
Séminaire DPhN, 2 février 2018
Collaborators: M. Fortin, F. Gulminelli, M. Hempel, M. Mancini, M. Marques, J.
Novak, C. Providencia, A. Raduta, S. Typel
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 1 / 30
http://luth.obspm.fr
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Outline
1 Introduction
2 What do we know?
3 Inhomogeneous matter
4 Homogeneous matter
5 Summary and Outlook
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 2 / 30
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A neutron star : a star made of neutrons....
1932, Landau (Phys. Z. Sowjetunion, 1, 285)Possibility of stars with a central density comparableto that of nuclei
1934, Baade and Zwicky (Phys. Rev. 45, 138)Prediction of the existence of neutron stars : Withall reserve we advance the view that supernovaerepresent the transition from ordinary stars intoneutron stars, which in their final stages consist ofextremely closed packed neutrons.
1939, Tolman, Oppenheimer, and VolkovGeneral relativistic neutron star models :M ≈ 1.5M� and r ∼ 10 km → density ∼ 0.1 fm−3
Walter Baade
Fritz Zwicky
Signal too weak to be observed → almost forgotten until 1967 :discovery of pulsars by Hewish and Bell
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 3 / 30
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Nowadays
Almost 3000 neutron stars havebeen observed as pulsars,among others Crab, Vela,Geminga, Hulse-Taylor, doublepulsar, . . .
Pulsars in many differentsystems
Estimate of age and surfacemagnetic field from measuredP, Ṗ (Rotating magnetic dipolemodel)
“Magnetars” have extremelyhigh magnetic fields(∼ 1015 G at the surface)
P -Ṗ diagram
[Becker et al., 1305.4842]
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 4 / 30
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Core-collapse supernovae : birth of acompact object
Supernova explosions observed since almost2000 yearsSN 1006, 1054, 1181, . . . : report by araband/or chinese astronomersTwo different classes
Thermonuclear explosionsI Type Ia Supernovae : ignition of
runaway nuclear fusion in a whitedwarf and subsequent explosion
Core-collapse supernovaeI All other types : gravitational collapse
and subsequent explosion of a massive(M >∼8-10M�) star at the end of itslife→ formation of a neutron star or astellar black hole
Crab nebula (Hubble telescope)
SN1987A
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Binary neutron stars
PSR 1913+16
First system observed in 1974 by Hulseand Taylor, nowadays ∼ 15 knownsystems
Extremely relativistic → preciseobservations allow for testing generalrelativity and determining neutron starmasses
Good source of gravitational waves→ change in orbital motion viagravitational wave emission : first(indirect) evidence for gravitationalwaves (Nobel prize for Hulse and Taylor1993)
August 17, 2017 : first detection of the coalescence of a binary neutron star (GW+ electromagnetic counterpart)No binary NS-BH system known
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 6 / 30
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Why studying dense and hot matter ?
What do we need as input from microphysics ?
Weak reaction rates (neutrino-matter interactions)
An equation of state
We want to understand :
Neutron starsI Several minutes( !) after their birth T ∼ n0/2 : nuclei disappear in favor of free nucleons
Binary neutron star mergers and neutron star black hole mergersI Close to merger matter is heated upI Very high densities reached in post-merger supermassive neutron star
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What is “hot and dense” ?
Large domains in density, temperature and electron fraction have to be covered
temperature 0 MeV ≤ T < 150 MeVbaryon number density 10−11 fm−3 < nB < 10 fm
−3
electron fraction 0 < Ye < 0.6
and matter composition changes dramatically throughout ! (Review MO et al RMP 2017)
Different regimes :
Very low densities and temperatures :I dilute gas of non-interacting nuclei → nuclear statistical equilibrium (NSE)
Intermediate densities and low temperatures :I gas of interacting nuclei surrounded by free nucleons → beyond NSE
High densities and temperatures :I nuclei dissolve
→ strongly interacting (homogeneous) hadronic matterI potentially transition to the quark gluon plasma
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 8 / 30
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Plan
1 Introduction
2 What do we know?
3 Inhomogeneous matter
4 Homogeneous matter
5 Summary and Outlook
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 9 / 30
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Constraints from nuclear physics
Energy per baryon of symmetricmatter
-20
-10
0
10
20
30
0 0.1 0.2 0.3 0.4 0.5
E/A
[M
eV]
nB [fm-3]
Nuclear masses (binding energies) for manynuclei close to stability
Extracting parameters of symmetric nuclearmatter around saturation (n0, EB ,K, J, L)
Data from heavy ion collisions (flowconstraint, meson production, . . .)
Data on nucleon-nucleon interaction fixingstartpoint of many-body calculations (dataon hyperonic interactions scarce)
Low density neutron matter : Monte-Carlosimulations and EFT approaches
Compact star matter not accessible in terrestriallaboratories (density, asymmetry) nor to ab-initocalculations !
Example : symmetry energy and slope(Lattimer & Lim 2013, MO et al 2017)
20
40
60
80
100
120
26 28 30 32 34 36 38 40 42
L (
MeV
)
J (MeV)
LSDD2
FSUgoldSka
IUFSUNL3
SFHoSFHxTM1TMA
TNTY
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Modelling a neutron star
Compactness Ξ =GM
Rc2∼ 0.2 → GR gravitation (maximum mass !)
Assumption of perfect fluid :Tµν = (p+ ε)uµuν + pgµν
cold matter + β-equilibrium→ EoS function of one parameter,p(nB), ε(nB) (or equivalent)
Simplest case : static equilibrium,spherical symmetry →TOV system→ global quantities such asgravitational/baryon mass and radius
Different EoS models (fromhttp ://compose.obspm.fr)
0
1
2
3
8 10 12 14 16
M (
M⊙
)
R (km)
Maximum mass is a GR effect, value given by the EoS
Numerical solutions exist for non-spherical cases (rotation, em field, . . .)Publicly available codes
I http ://www.lorene.obspm.frI http ://www.gravity.phys.uwm.edu/rns/
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Constraints from observations
Observations Quantities detected Dense matter properties
Orbital parametersin binary systems
Neutron star massesEquation of state (EoS),high densities
GW from binary systems Tidal deformability Compactness, EoS
Pulsar timing GlitchesEvidence for superfluidcomponent
X-ray observations Surface temperatureHeat transport/neutrinoemission, superfluidity
RadiiEoS, also low and interme-diate densities (crust)
Pulsar timing NS rotation frequencies EoS via mass-shedding limit
GWs OscillationsEigenmodes (EoS, crustproperties)
QPO Radii EoSAsterosismology Eigenmodes
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Neutron star observations1. Neutron star masses
Observed masses in binary systems(NS-NS, NS-WD, X-ray binaries) withmost precise measurements from doubleneutron star systems.
Two precise mass measurements inNS-WD binaries
I PSR J1614-2230 :M = 1.928± 0.017M� (Fonseca et al 2016)
I PSR J0348+0432 :M = 2.01± 0.04M� (Antoniadis et al 2013)
Given EoS ⇔ maximum massAdditional particles add d.o.f.→ softening of the EoS→ lower maximum mass→ constraint on core composition
0.6 0.8 1 1.2 1.4 1.6 1.8 2 2.2
MG [Msol]
Companion
J1829+2456
Companion
J1811-1736
Companion
J1906+0746
Companion
B1534+12
Companion
B1913+16
Companion
B2127+11C
J0737-3039B
J0737-3039A
Companion
J1756-2251
Companion
J1807-2500B
Companion
J0453+1559
J1141-6545
J2043+1711
J1918-0642
J1910-5958A
J2222-0137
J1802-2124
J1713+0747
J2234+0611
J0437-4715
J1738+0333
J1909-3744
J0621+1002
B1855+09
J0751+1807
J1012+5307
J1946+3417
J1614-2230
J0348+0432
J0337+1715
J1903+0327
NS-NS binaries
WD-NS binaries
Triple systems
(update from Lattimer & Prakash, 1012.3208)
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Neutron star observations2. Radius estimates from X-ray observations
Radii from different typesof objects, but modeldependent :
I Atmosphere modelling(much recent progress)
I Interstellar absorption(X-ray observations)
I Distance, magneticfields, rotation, . . .
Many discussions
Consensus : radius of a fiducialM = 1.4M� star 10-15 km
2σ error bars, radii at M = 1.4M�
8 9 10 11 12 13 14 15 16R1.4 [km]
BNS+QXT
QXT-1
BNS-2
BNS-1
RP-MSP New mass
He atmosphere?
(courtesy M. Fortin, CAMK)
See also new determination by J. Nättilä et al R(1.9M�) = 12.4 ± 0.4km
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 14 / 30
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Neutron star observations3. Pulsar timing
Measurements of rotational frequencyI f = 716 Hz (PSR J1748-2446ad) (Hessels et al, Science
2006)
Theory : Kepler frequency (Haensel et al. A&A 2009)fK = 1008 Hz (M/M�)
1/2 (R/10km)−3/2
APR
FPS
DH
BGN1H1
GMGS
BM165
SQM3 SQM1
(Courtesy M. Fortin, CAMK)
NS radii, dependence oncrust-core transition
12 13 14 15R (km)
0.5
1.0
1.5
2.0
2.5
M (M
⊙)
1.0 1.5 2.0 2.5M (M⊙)
0.5
1.0
1.5
2.0
2.5
3.0
lcr (k
m)
Unifiednm=0.01nm=ncnm=ntnm=n0nm=0.5n0−n0nm=0.1n0−nt
(Fortin et al. 2016)
→ A measured frequency of1.4 kHz would constrainR1.4 < 9.5 km !Remark of cautionRadii are sensitive torotation and to treatment ofcrust-core transition !
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GW from binary NS mergers
GW170817 : first detection of a NS-NS merger with LIGO/Virgo detectors
Information on EoS from different phasesI Inspiral → masses of objectsI Late inspiral → tidal deformability Λ̃
depends on matter properties(Read et al, Faber & Rasio, Hinderer et al . . .)
GW170817
Λ̃ < 800 (90% confidence level)(low spin prior) (Abbott et al 2017)
I Post merger oscillations → peak frequencystrongly correlated with NS radius(Bauswein et al, Sekiguchi et al, . . .)
12 13 14 150.7
0.8
0.9
1
1.1
1.2
R1.6
[km]
f pe
ak/M
tot [k
Hz/M
su
n]
2.4 Msun
2.7 Msun
3.0 Msun
(Bauswein et al 2016)
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Tidal deformabilityConstraints on the EoS
Tidal deformability Λ̃depends on matterproperties
Λ̃(Mchirp , q,EoS)
∼ 5% uncertainty fromcrust treatmentpreliminary !
. 10% uncertainty fromthermal effects preliminary !
Tidal deformability for different EoS, q = M1/M2
0
400
800
1200
1600
0.6 0.7 0.8 0.9 1
Λ
q
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Tidal deformabilityMass-radius relations
Some EoS (giving less compact NSs) excluded by limit on Λ̃
Additional (modeldependent) constraintsfrom relation with EMobservations
I Mtot + no prompt BHcollapse (Bauswein et al 2017)
I Mtot + estimate ofenergy loss to ejecta(Margalit& Metzger 2017)
I Ejecta masses +composition (Shibata et al 2017)
I Λ̃ & 450 from ejectamasses(Radice et al 2017)
Different EoS models compatible with Mmax > 1.97M�
0
1
2
3
8 10 12 14 16
M (
M⊙
)
R (km)
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Tidal deformabilityMass-radius relations
Some EoS (giving less compact NSs) excluded by limit on Λ̃
Additional (modeldependent) constraintsfrom relation with EMobservations
I Mtot + no prompt BHcollapse (Bauswein et al 2017)
I Mtot + estimate ofenergy loss to ejecta(Margalit& Metzger 2017)
I Ejecta masses +composition (Shibata et al 2017)
I Λ̃ & 450 from ejectamasses(Radice et al 2017)
Different EoS models compatible with Mmax > 1.97M� and Λ̃
0
1
2
3
8 10 12 14 16
M (
M⊙
)
R (km)
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 19 / 30
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Tidal deformabilityMass-radius relations
Some EoS (giving less compact NSs) excluded by limit on Λ̃
Additional (modeldependent) constraintsfrom relation with EMobservations
I Mtot + no prompt BHcollapse (Bauswein et al 2017)
I Mtot + estimate ofenergy loss to ejecta(Margalit& Metzger 2017)
I Ejecta masses +composition (Shibata et al 2017)
I Λ̃ & 450 from ejectamasses(Radice et al 2017)
Different EoS models compatible with Mmax , Λ̃ and EM constraints
0
1
2
3
8 10 12 14 16
M (
M⊙
)
R (km)
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 20 / 30
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Plan
1 Introduction
2 What do we know?
3 Inhomogeneous matter
4 Homogeneous matter
5 Summary and Outlook
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 21 / 30
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Clustered matterNuclear abundances important for composition of (proto-)neutron star crust,nucleosynthesis and CCSN matterModelling does not only depend on the interaction chosen :
Theoretical description of inhomogeneous system (interplay of Coulomb andstrong interaction, surface effects, . . .)
Binding energies of (neutron rich) nuclei
Treatment of excited states
Transition to homogeneous matter (stellar matter is electrically neutral !)
Nuclear abundances within different models (same thermodynamic conditions, gas density negligible) (MO et al 2017)
0
10
20
30
40
50
60
70
Z
0 10 20 30 40 50 60 70 80 90 100 110 120
N
log10(XN,Z)
-10-8-6-4-20
a) SMSM
nB = 5.69 10-4
fm-3
T = 1.60 MeV
Ye = 0.35
0
10
20
30
40
50
60
70
Z
0 10 20 30 40 50 60 70 80 90 100 110 120
N
log10(XN,Z)
-10-8-6-4-20
b) Gulminelli and Raduta (2015)
0
10
20
30
40
50
60
70
Z0 10 20 30 40 50 60 70 80 90 100 110 120
N
HS(DD2)
STOS
LS220
log10(XN,Z)
-10-8-6-4-20
c) HS(DD2)
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Reaction rates
Overall reaction rates : matter composition + individual ratesI Homogeneous matter : calculate individual rates in hot and dense medium
→ collective responseI Clustered matter : rates on nuclei far from stability (up to now essentially shell
model)
Different (weak) interactionrates are extremely important !Neutrino emission, electroncapture, . . .
And very sensitive to thedifferent ingredients
I Example : influence ofnuclear masses for nuclei withneutron numbers betweenN = 50 and N = 82→ up to 30% change inoverall EC rate
Nuclear masses and EC rates
-0.2
-0.1
-0
0.1
0.2
0.3
10-6
10-5
10-4
10-3
∆Z=2 25Msun
∆Z=2 15Msun
∆Z=5 25Msun
∆Z=5 15Msun
∆Z=10 25Msun
∆Z=10 15Msun
nB (fm
-3)
< λ
ecm
> / <
λec >
- 1
0
0.5
10-4
10-3
(Raduta et al 2016)
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 23 / 30
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Plan
1 Introduction
2 What do we know?
3 Inhomogeneous matter
4 Homogeneous matter
5 Summary and Outlook
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 24 / 30
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Particle content : there is nothing exotic !On the hadronic side, non-nucleonic degrees of freedom such as hyperons andmesons are well known and studied experimentally
PionsI discovered in 1947 in cosmic raysI by far lightest hadrons (mπ ∼ 140 MeV)I prominent member of the nuclear
interaction
Kaons, strange cousins of the pions(mK ∼ 500 MeV)Hyperons
I first hypernuclei by Daniesz and Pniewskiin 1952
I studied in scattering experiments
Nuclear resonances (∆, . . .)
��
���
TTTTT
TT
TTT
�����
r
r
r
rr rrr
n p
Ξ− Ξ0
Σ+Σ−Σ0
Λ
s = 0
s = −1
s = −2
1
Charged leptons other than e± :muons (mµ = 105 MeV) and tauons (mτ = 1.8 GeV)
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 25 / 30
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Where do hyperons appear ?
regions with an overall hyperon fraction > 10−4 (Marques et al 2017)
10-6
10-4 0.01 1
YQ = 0.1
nB (fm-3
)
10-4 0.01 1
YQ = 0.3
nB (fm-3
)
10-4 0.01 1
0
10
20
30
40
50
60
YQ = 0.5
nB (fm-3
)
T (MeV)
Low charge fraction favors hyperons
Hyperons change matter properties at high densities and temperatures (PNS,NS, merger remnant)
Not only EoS is important, but transport properties, too (e.g. hyperonicDURCA for NS cooling (Fortin et al 2016, Raduta et al 2017))
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 26 / 30
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I-Q relation with hyperons
I-Love-Q relations : universal (independent of the EoS) relations betweenmoment of inertia, quadrupole moment and deformability for cold slowlyrotating NSs (Yagi & Yunes 2013)
Rotation frequencydependent
EoS independent ifhyperons are included, too
Relations becometemperature dependent asfar as thermal effects startto modify the EoS, i.e. forearly PNSs and mergerremnants (Marques et al 2017)
0
0.02
0.04
0.06
0.08
0.1
0.12
1 2 3 4 5 6 7
∆I-- / I-- fi
t
Q--
sB = 0
sB = 1
sB = 2
sB = 4
What about thermal effects for deformability in late inspiral of a BNS merger ?
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 27 / 30
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Plan
1 Introduction
2 What do we know?
3 Inhomogeneous matter
4 Homogeneous matter
5 Summary and Outlook
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 28 / 30
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Summary
We need to know matter properties (EoS and reaction rates) in regions notaccessible to experiments !
Much progress in recent years on general purpose EoSI Inhomogeneous matter : statistical models with entire nuclear distributionI Homogeneous matter : non-nucleonic degrees of freedom included consistent
with contraints
Not only nuclear matter parameters are important ! Modelling ofinhomogeneous matter influences CCSN (weak reactions, . . .) and NS (radii,. . .)Detailed matter composition important for dynamical evolution (transportcoefficients, . . .), too
ComPOSE data base (T. Klähn, M. Mancini, MO, S. Typel)
Many EoS models available in tabular form fromhttps://compose.obspm.fr
Compose software for interpolation, calculation of thermodynamic quantities(sound speed, adiabtic index, . . .) and extracting compositional information
Coupled to Lorene for cold neutron stars (J. Novak)
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 29 / 30
https://compose.obspm.fr
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Outlook
1. GW detectorsI new NS mergers events → deformabilityI post-merger phase probably out of range for present detectorsI kilonova and r-process nucleosynthesis needs detailed matter composition
Moment of inertia (Livre blanc SKA France)
1
1.2
1.4
1.6
1.8
2
10 12 14 16 18
I (1
038 k
g m
2)
Req (km)
HS(DD2)STOSSLY9
SLY230a
2. Radiotelescope SKAI Many new pulsars with precise mass
determinationsI Radius via moment of inertia ?I Fast rotating objects, constraint via
Kepler frequency ?
3. NICER, . . .I Soft X-ray timing and spectroscopyI Main goal : radii to < 5% precision
And. . .
Most general purpose models not compatible with constraints ! Improvereliability (modern interactions compatible with constraints, coupling with abinitio methods, nuclear data for neutron rich nuclei, . . .)
EoS and reaction rates should be treated consistently
Micaela Oertel (LUTH) Neutron stars Saclay, February 2, 2018 30 / 30
IntroductionWhat do we know?Inhomogeneous matterHomogeneous matterSummary and Outlook