Broadband Emission of Magnetar Wind Nebulae Shuta Tanaka ICRR, The Univ. of Tokyo 29, Oct., 2015, TeVPA 2015 @ Kashiwa 1
Contents 1. Introduction 2. Spectral Evolution of young PWNe 3. A Spectral Model of MWNe 2
Pulsars in PWNe Crab (Chandra) l Pulse lumi. ~ a few % x L spin Most of L spin releases as pulsar wind! l ~ 50 of 2000 pulsars have observable PWNe. (L spin > 10 36 erg/s). P - P diagram 10 36 erg/s 3C58 (Chandra + VLA) L spin = IΩ Ω l Bow-shock PWNe are around MSP of L spin < 10 36 erg/s How about other classes of pulsar? 3
Magnetar Wind Nebulae? PWNe are found around High-B radio pulsar. One of young TeV PWN around high-b radio pulsar Observed extended emission may be dust-scattering halo for 1E1547.0 (e.g., Olausen+11) 4
Magnetar Wind l Magnetars have large surface B-field ( > 3x10 14 G) and some have L x > L spin B-powered emission?. l Magnetars have P Wind loss Ṁ wind R LC L wind c = L spin for RPP L X [erg/s] Angular momentum loss!! L x = L spin Photon loss Ṁ ph (R NS + ) L ph B c NS [G] L spin [erg/s] Even L x > L spin, wind angular momentum loss would dominate for magnetars from this simple estimate. 5
What We Learn from MWN l PWN spectrum tells us spin evolution of pulsar => How about MWN? l Does a millisecond magnetar exist? l Difference of wind properties from rotation powered pulsar (RPP). l magnetization: σ? pair multiplicity: κ? l Testing fall-back disk model of magnetar. l Disk model does not produce MWN because no wind flows out from magnetosphere. 6
Contents 1. Introduction 2. Spectral Evolution of young PWNe 3. A Spectral Model of MWNe 7
Spectral Evolution Model l PWN as uniform sphere l Constant rate of Expansion applicable to young (<<10kyr) PWN R PWN = V PWN x t e, B l Introduce a parameter η satisfies (no hadron) L spin = L e± + L B = (1-η) L spin + ηl spin l B-field evolution as energy conservation t 0 L B (t ')dt ' Total magnetic energy injected = 4π 3 R 3 B 2 (t) PWN 8π Magnetic energy inside PWN 8
Spectral Evolution Model l Injection of non-thermal e ± plasma (broken power-law) (1-η)L spin = L e± = Q inj γmc 2 dγ logn γ -p1 γ L -p2 e γ min γ b γ max logγ l Evolution of energy distribution of e ± t Radiative & adiabatic cooling N( γ, t) +! γ ( γ, t) N( γ, t) = Qinj( γ, t) γ Injection from pulsar Calculate Syn. & IC spectral evolution 9
Crab Most observed & studied object (one parameterη) Spectral Evolution of the Crab Nebula l Observed flux@1kyr η~ 5 x 10-3 (85μG) synchrotron Inverse Compton l SSC dominates in γ-rays. S.T.& radio[%/yr] Opt.[%/yr] Model -0.16-0.24 Obs. -0.17-0.55 l Flux evolution is consistent. Model can reproduce observations 10
G292.0+1.8 G54.1+0.3 Young PWNe Chandra Park+07 G21.5-0.9 5.2. G21.5-0.9 Applicable to other young69 PWNe < 10kyr G11.2-0.3 10-9 Chandra G0.9+0.1 νfν[ergs/cm2/sec] 10-10 G21.5-0.9 G21.5-0.9 (Model 2) 3C58-11 10 Kaspi+01 10-12 B0540-69.3 SYN IC/CMB IC/IR IC/OPT SSC Total -13 10 10-14 10-15 1010 Porquet+03 1015 1020 ν[hz] 1025 S.T.&Takahara11 Park+10 S.T.&Takahara 13 Figure 5.10: Model spectrum of G21.5-0.9 at tage = 1.0kyr for model 2, where we ignore the One-zone model is enough to get distribution (right panel) G21.5-0.9 with the use of the same parameters as in Figure 5.8. l ofmean B-field inside PWN. While the synchrotron flux decreases with time, the inverse Compton flux increases. This l spin-evolution of central pulsars feature is due to the decrease of the magnetic field strength with time and to the increase of the particle number as seen in the right panel. Because the energy injection from the pulsar will l predictions to γ-ray flux continue till the time t τ = 3.9kyr, the particles in the PWN increase (see the discussion in -3 for lin the η~10 allof theexcept Section 5.1.2 in details). left panel, we can see the evolution synchrotron cooling for Kes 75 infrared observation (Gallant & Tuffs, 1998) in the spectral fitting. Used and obtained parameters Kes 75 3C58 are tabulated in Table 5.2. 3C58 (Chandra + VLA) Kumar & Safi-Harb08 0 break frequency νc (t) (from 7 1014 Hz at 300 yr to 1017 Hz at 10 kyr) as is predicted by 11
Contents 1. Introduction 2. Spectral Evolution of young PWNe 3. A Spectral Model of MWNe 12
Simplified One-zone Model l Spin-down evolution l Expansion of MWN l B-field inside MWN l e ± injection logq γ -p1 L e γ -p2 γ min γ b γ max logγ t c = n 1 (t age + 0 ) 2 L spin (t) L spin,now t t age R(t) =R now B(t) =B now L e ±(t) n(,t)+ t t t age R L spin (t) e ± energy distribution e, B n(,t) =Q(,t) adiabatic + synchrotron + inverse Compton coolings t t age n 1 2-1.2 < α R < -1.0, -5 < α B < -3 from past studies B 2 t age n+1 n 1 13
Kes 75 Young RPP with magnetar-like activity L spin = 8x10 36 erg/s, B NS = 4.8x10 13 G, τ c = 0.7kyr model 1 model 2 d 6kpc 10.6kpc τ 0 0.2kyr 3yr E rot 1.5x10 48 erg 2.1x10 50 erg B PWN 20μG 24μG η 5x10-5 6x10-6 14
νf ν [ergs/cm 2 /sec] 10-11 10-12 10-13 10-14 10-15 10-16 AXP 1E1547.0-5408 Assuming (u IR, u opt ) = (1.0, 2.0) [ev/cc] L spin = 1x10 35 erg, B NS = 2.2x10 14 G, τ c = 1.4kyr AXP 1E1547.0-5408 (Model 1) (B 1, γ b, p 2 ) = (3µG, 10 6, -2.6) SYN IC/CMB IC/IR IC/OPT TOTAL XMM-Newton 10 10 10 15 10 20 10 25 ν[hz] CTA l Results are insensitive to α B & α R. l X-ray upper limit gives constraints B now < 25μG. l CTA would detect MWN around 1E1547.0 when B now < 3 μg & u IR > 1.0eV/cc l η<<10-3 is consistent with Kes 75. l P 0 < 100msec is presumed. νf ν [ergs/cm 2 /sec] 10-11 10-12 10-13 10-14 10-15 10-16 AXP 1E1547.0-5408 (Model 2) (B 1, γ b, p 2 ) = (25µG, 10 4, -2.6) SYN IC/CMB IC/IR IC/OPT TOTAL XMM-Newton 10 10 10 15 10 20 10 25 ν[hz] CTA 15
Summary l Pulsar & Pulsar Wind Nebula l Most of L spin lost by pulsar wind l σ- & κ- problem about physics of pulsar magnetosphere l One-zone model of PWNe tells us spin-evolution of central pulsars l A Spectral model of MWNe l Angular momentum loss of magnetar => magnetar wind? l Deep obs. of 1E1547.0 by CTA can examine the hypothesis. l η<<10-3 is consistent with High-B radio pulsars. l Constraint on P 0 is not strong < 100msec l Detections of MWN in multi-wavelength are required to constrain magnetar evolution properties. 16