neutron stars: probing ultra-dense and hot...

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

Outline

1 Introduction

2 What do we know?

3 Inhomogeneous matter

4 Homogeneous matter

5 Summary and Outlook


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


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]


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


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


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


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


Plan

1 Introduction

2 What do we know?





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


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/


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


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)


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


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 !


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)


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


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)









Different EoS models compatible with Mmax > 1.97M� and Λ̃

0

1

2

3

8 10 12 14 16

M (

M⊙

)

R (km)









Different EoS models compatible with Mmax , Λ̃ and EM constraints

0

1

2

3

8 10 12 14 16

M (

M⊙

)

R (km)


Plan

1 Introduction

2 What do we know?





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)


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)


Plan

1 Introduction

2 What do we know?





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)


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))


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 ?


Plan

1 Introduction

2 What do we know?





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)


https://compose.obspm.fr

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


IntroductionWhat do we know?Inhomogeneous matterHomogeneous matterSummary and Outlook

neutron stars: probing ultra-dense and hot...

Documents