3D Dynamical Modelingof the Gamma-ray Binary
LS 5039

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

1	  Outlinel Introduction to the VHE gamma-raybinariesl LS 5039: Previous resultsl LS 5039: 3D SPH simulationsl Concluding remarks

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Introduction to the VHE
Gamma-ray Binaries

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VHE gamma-ray binaries (1)lOnly 5 systems, all of which consist of anOB star and a compact object.
Ø3 systems with a Be starØ

2 systems with a main-seq. O starl

Nature of the compact object establishedonly for one system (PSR B1259-63).

l

Two competing scenarios for other systems:
Pulsar wind (PW) scenario vs.
Microquasar (MQ) scenario
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Be Stars
H"αNon-supergiant B
stars withemission lines
Viscous decretion
diskCTA	  Japan	  WS	  2013	  (September	  3-­‐4)	
Courtesy of Stan Owocki

Hydrogen"spectrum"
Intensity"

H"βλ0 Wavelength"5	

VHE gamma-ray binaries (2)System	Scenario	
Optical Porbstar	 (d)	e	PSR B1259-63	

PW	Be	1237	 0.87	

LS I +61 303	HESSJ0632+057	
LS 5039	
1FGL
J1018.6-5856	

MQ(?)	
?	

Be	Be	PW(?)	
?	

O	O	

26.5	 0.54	
321 0.83	(315)	
3.9	 0.35	
16.6	 low	CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Our motivationsl Why do only massive binaries show VHE

gamma-rays?
l Does the PW scenario work for all VHE
gamma-ray binaries? Or, does the MQ
scenario work for some systems?
l In systems with Be stars, what’s the roleof the Be disk for high energy emission?

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

7	Study of VHE gamma-ray binaries
3D dynamical
Is a new frontier ofmassivehelps	
starmodeling
research.	  The	  current	  goals	  are:	  	  llto	  understand	  the	  ature	  of	  ofindividual	  tounderstandthennaturesystems,	  including	  individual
systems, including
physics	  ooff	  interacGons	  Øphysics
interactions
origin	  ooff	  hhighigh	  eenergy
nergy	  emission	  Øoriginemissionto	  eestablish
stablish	  aa	  classiﬁcaGon/uniﬁcaGon	  toclassification/
scheme	  of	  these	  systems	  unification
schemeof these systems
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Shape of the colliding-wind interaction front
without orbital motion
Momentum balance determines the bow
shock structure (e.g., Canto et al. 1996)

M 2 v2
η≡M 1v1
0.03

0.25
0.5

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

9	  Interaction in a Be-star system (1)PSR B1259-63!

LS 2883

radio pulsar!

O9.5e!

Porb  3.4 yr,
e = 0.87
CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	  10	  Interaction in a Be-star system: an example
Widely-accepted picture of the PW-Be diskinteraction in PSR B1259-63	
Is thisrealistic?	

Porb  3.4 yr,
e = 0.87(NASA)	
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complicatedρ0 = 10 −9 g cm −3

Pulsar passes throughBe disktermination of PW
over a large solid
angle
Efficientconversion of PW
power to nonthermal emission

Density on orbital plane!

(Takata+ 2012)	

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Structure of shocked winds is timedependent and 3D! (Takata+ 2012)

Density on
orbital plane	CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

PW pressure on
orbital plane	
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High energy emission from the shocked
pulsar wind (2) (Takata+ 2012)>300 GeV Light Curves	

Double peak	ρ 0 = 10 −9 g cm −3

ρ 0 = 10 −10 g cm −3
ρ 0 = 10 −11 g cm −3-20	 -10	 0	 10	 20	 30	 40	CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	
Days	

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LS 5039: Previous results!

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Observed features
O6.5V + compact object: Porb=3.906d, e=0.35(Casares 2013)(Aharonian et al. 2006)
CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Fermi & HESS
Observationsll

Fermi 0.1-10 GeV
peaks near sup.
conj. (Abdo et al.
2006)
HESS 1-40 TeV
peaks near inf. conj.
(Aharonian et al. 2006)l

X-ray lightcurve
similar to that in
TeV (Takahashi et al.
2009)

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Observed features that favor PW scenario(Marc Ribo@Variable Galactic Gamma-ray Sources)
l

l

l

VLBA map regularly varies with orbital
phase, and is compatible with PWscenario.
All emission in radio, X-rays, GeV
gamma-rays, and TeV gamma-rays
behave similarly. No MQs do that.
LS 5039 shows no state change. MQs
show jets at some particular states.
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VLBI map + a simple model
inclination angle 60-75 deg.
(Moldon et al. 2012)	  	Emitting region adopted by Moldon et al. (2012)	

But, the PW cone can’t be this thin.	
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HE/VHE Gamma-ray emission calculation
via relativistic 2D simulations
Zabalza et al. (2012)	

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Thermal X-ray flux/spectrumll

So far, no significant thermal X-rays havebeen detected for any VHE binaries.
The upper limit of EW(Fe line) is 40eV(Takahashi et al. 2009)ll

But, thermal X-rays could be detected by
future missions, such as ASTRO-H.
Even if it is too weak to be detected, it canbe used to constrain the spin down
luminosity of the pulsar (Zabalza et al. 2011).
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Semi-analytical shock structuere modelby Zabalza et al. (2011)
Lsd = 3 × 10 35 erg/s
Lsd = 3 × 10 erg/s
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Lsd = 3 × 10 37 erg/s

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Thermal X-ray spectra(Zabalza et al. 2011)

Lsd = 3 × 10 37 erg/s
Lsd = 3 × 10 36 erg/s
Lsd = 3 × 10 35 erg/s

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Upper limit of pulsar spin down luminosity(Zabalza et al. 2011)

Lsd ≤ 6 × 10 erg/s
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CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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LS 5039: 3D SPH
simulations

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Numerical method: Hydro simulationsll

3D Smoothed Particle Hydrodynamics code
βStellar wind: vw = v∞ (1 − R / r) with

β =1l

l

Pulsar wind: relativistic pulsar windmodeled by a high-velocity, nonrelativistic wind with the same
momentum flux
Optically-thin radiative cooling
CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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SPH (Smoothed Particle Hydrodynamics)

Density

A particle method that divides fluid
into a set of discrete "fluidelements” (=particles).

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

x27	

Numerical method: radiative transferobserver
computation of theradiative transfer fora series of 1D rays
for every energy E	

jE = ne ni Λ(E,T )κ (E)systemCTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Stellar, wind, and orbital parameters for
LS 5039 accretion simulations
Primary
SecondarySpectral TypeO6.5V
pulsar
22.9Msun
1.4MsunMass--9.3Rsun
RadiusVinf
2,440 km/s 12,200 km/s--Twind39,000 K--Mdot
2.5x10-7Msun/yr
3.9060 days
Porb
Eccentricity
0.35CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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3D structure of the interaction surfacell

PW carvesout stellarwind.
PW region
has similardimension inr and z

Lsd = 5 × 10 36 erg/s
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large-scale structure ofinteraction surface

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Effect of PW luminosity
Lsd = 5 × 10 erg/s
37

Lsd = 5 × 10 erg/s
35

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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Simulated thermal X-ray spectra

CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

33	  Simulated thermal X-ray light curve
Lsd = 5 × 10 36 erg/sCTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

Thermal X-ray
flux is higher at
INFC than at
SUPC becauseof lower
absorption andhigher stellar
wind velocity at
INFC.	
34	Concluding remarksl

Simulated thermal X-ray flux is about oneorder of magnitude lower than theanalytical one.
Higher upper limit of pulsar spin down
luminosity by ~10.

l

3D numerical simulations will provide a
powerful tool to understand the nature ofindividual systems and establish aclassification/unification scheme of VHE
gamma-ray binaries.
CTA	  Japan	  WS	  2013	  (September	  3-­‐4)	

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