Accretion Models of Gamma-Ray Bursts

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PhysicsofFlaresfromBlackHoleMicroquasarV4641SgrintheRadio,OpticalandX-rays

Soon-WookKim12andShinMineshige3

†

RadioAstronomyDivision,KoreaAstronomyObservatory,Korea†YukawaInstituteforTheoreticalPhysics,KyotoUniversity,Japan

Abstract.Wepresentadiskinstabilitymodelforunusuallyshort,butintenseandchaotic,ﬂaresofafewdaysobservedinablackholemicroquasarV4641Sgr=SAXJ1819.3-2525,amultiplecomponentrelativisticjetsource.Toreproducetheobservedshortdurationoftheﬂare,aring-like,truncatedKepleriandiskisrequired.Wediscusscausesofsuchatruncatedaccretiondiskbasedonthediskinstabilitymodel,andimplicationtoshort-durationﬂareeventsobservedinothermicroquasarsandfastX-raytransients.

BLACKHOLEMICROQUASARV4641SAGITTARII=SAXJ1819.3-2525

BlackholemicroquasarV4641SgrwasdiscoveredinFebruary1999inquiescencebyBeppoSAXandRXTE/AllSkyMonitor,anddesignatedasSAXJ1819.3-2525[1][2].OnSeptember1999itsbrightandsuper-Eddington,butunusuallybriefoutburstwasobservedintheradio,opticalandX-raysupto100keV[3][4][5].Ac-companiedsuperluminalradiojetidentiﬁesthesourceasamicroquasar[3].OpticalspectroscopicobservationsshowthatV4641Sgrisahighmassblackholebinarywiththeblackholeandcompanionmassesof9.6and6.5M,respectively,withanorbitalperiodof2.81days[6].InFigure1,wepresentthe1999FebruarySeptemberX-rayandopticallightcurvesofV4641Sgr.Inthetopandmiddlepanels,X-rayobservationsaredisplayed:publiclyavailabledataofRXTE/PCA(ﬁlledandopentriangles),RXTE/ASM4()andBeppoSAX(opensquare)[5].Theopticallightcurveisalsopresented(adottedline)inthetoppanelforacomparison.Inthebot-tompanel,thevisualobservations(opencircle)adoptedfromVSNET5andfrom[7],andCCDdata(ﬁlledcircle)[4]arepresented.Intheoptical,theprecursorﬂaringactivitiesforaoftwoandahalfyears.Inthispaper,wefocusonlyonthe1999outburst.

脉冲星的磁层模型与辐射特性分析

脉冲星的磁层模型与辐射特性分析脉冲星是宇宙中一种高度致密的天体，由恒星演化过程中碳氧燃烧结束时产生的中子星引力塌缩形成。

它们具有极强的磁场，导致它们以极快的自转速度旋转，并产生规律的脉冲射电、X射线、γ射线等辐射。

脉冲星的磁层模型可以通过研究它们的辐射特性来进行分析。

其中一种常见的模型是磁力线束模型。

根据这个模型，脉冲星的极强磁场将导致电子在磁力线上进行快速加速和减速，形成辐射束。

当辐射束指向地球时，我们就能观测到脉冲星的脉冲信号。

磁力线束模型提供了解释脉冲星辐射特性的重要线索。

通过分析脉冲星的光度曲线，我们可以确定其自转周期和脉冲宽度。

同时，脉冲星的脉冲轮廓也可以提供关于磁力线束几何形状的信息。

不同形状的束流可以解释不同的脉冲特征，比如双峰、多峰、宽度变化等。

此外，研究脉冲星辐射特性的另一种方法是通过频谱分析。

脉冲星的射电辐射通常呈现出不同频率上的特征波峰。

这些波峰在不同频率下的位置和强度变化提供了关于磁场结构和加速机制的信息。

通过测量不同频段的射电脉冲信号，我们可以揭示脉冲星磁层中电子加速和辐射的机制。

此外，X射线和γ射线也是研究脉冲星辐射特性的重要手段。

由于强磁场和高自转速度产生的极高加速度，脉冲星可以产生高能射线。

这些射线源自于磁层中的高能电子和正电子相互作用和湮灭过程。

通过测量不同能段的X射线和γ射线辐射，我们可以了解脉冲星活动的高能过程、能谱特性等。

磁层模型与辐射特性分析不仅有助于我们理解脉冲星的内部结构和物理过程，还为我们研究宇宙物理学提供了一个重要的实验场所。

脉冲星作为极端物理现象的代表，其磁场、自转等特性对于研究宇宙中的磁场演化、星际介质的物理性质等课题具有重要意义。

此外，脉冲星辐射特性的研究还对于精确测量宇宙中的距离、探测暗物质等具有潜在应用价值。

总之，脉冲星的磁层模型与辐射特性分析是一项复杂而有深度的研究课题。

通过研究脉冲星的辐射特性，我们可以了解到脉冲星的内部结构、磁场形态以及加速辐射机制等重要信息。

Extragaactic X-ray Binary Stars Beacons for Stellar- and ：河外星系的X射线双星的恒星和信标

Best spectra show non-solar abundance ratios.
Metallicity of Hot Gas in Ellipticals
MOS+PN for NGC4472, NGC4649, and NGC1399
O/Fe ~ 0.3 - 0.4 Mg/Fe ~ 1.2 - 1.5 Ni/Fe ~ 3 - 7
Stellar Source for Soft Component?
M31 (diffuse) LX (0.5-2 keV): 1.7 x 1038 ergs s-1 LX (2-7 keV): 5.3 x 1037 ergs s-1
LX(0.5-2 keV)/ LX(2-7 keV) = 3.2
Total 0.5 - 2.0 keV flux Source: 85% Background: 15%
Source 0.5 - 2.0 keV flux Gas: 85% LMXBs:15%
Source of Low Metallicity Gas
How are both LX/Lopt and low metallicity achieved? One solution: ongoing accretion of pristine gas surrounding galaxies dilutes to subsolar metallicities
Cause of LX vs. Lopt Dispersion
Environmental? ram pressure stripping: LX - cluster environment ICM pressure confinement: LX “stifling” Internal? - variation in Type Ia supernovae-driven winds - variation in depth of dark matter gravitational potential Low LX/Lopt in low density environments. High and low LX/Lopt in high density environments. (e.g., Brown & Bregman 2000)

费米选耀变体序列研究

费米选耀变体序列研究王雪品;王兴华;丁楠;李丙郎;李斯;伍林【摘要】收集了129个费米选耀变体的多波段准同时性观测数据,利用对数抛物线拟合其能谱分布(SED),获得同步辐射峰值频率(vs)、同步辐射峰值光度(Ls)及其他相关参数.分析结果表明:Ls与vs之间有明显的负相关性,支持了耀变体演化遵循从FSRQs到LBLs到HBLs的顺序;星系中心存在小质量黑洞导致Ls-vs图中呈现低vs且低Ls的耀变体.【期刊名称】《云南师范大学学报（自然科学版）》【年(卷),期】2016(036)001【总页数】6页(P1-6)【关键词】同步辐射;中心黑洞质量;耀变体序列;相关分析【作者】王雪品;王兴华;丁楠;李丙郎;李斯;伍林【作者单位】云南师范大学物理与电子信息学院,云南昆明650500;云南师范大学物理与电子信息学院,云南昆明650500;云南师范大学物理与电子信息学院,云南昆明650500;云南师范大学物理与电子信息学院,云南昆明650500;云南师范大学物理与电子信息学院,云南昆明650500;云南师范大学物理与电子信息学院,云南昆明650500【正文语种】中文【中图分类】P14耀变体是一类具有高光度、快速光变、高偏振、非热连续辐射、视超速运动等特征的射电噪活动星系核(AGNs).耀变体的辐射光子能量范围从射电到伽马射线，其多波段能谱分布(SED)呈现双峰结构[1-2].低能峰位于从近红外到X射线能段，一般认为其产生于喷流内极端相对论性电子的同步辐射[3-4]；高能峰位于从MeV到GeV γ射线能段，产生的原因目前还存在很大争议.轻子起源模型认为：高能峰产生于极端相对论电子逆康普顿散射自身发射的同步辐射光子(SSC)，或产生于极端相对论电子逆康普顿散射喷流外部的低能光子(EC)[5-6].耀变体分为平谱射电类星体(FSRQs)和蝎虎天体(BL Lacs)，FSRQs有强的发射线，而BL Lacs只有很弱或没有发射线(等值宽度<5 Å)[7].Fossati等人提出一个FSRQs和BL Lacs统一的耀变体能谱演化序列[8]，该演化序列表明：辐射功率越大的源，其SED中同步辐射的峰值频率和逆康普顿散射的峰值频率就越低，而康普顿主导度(CD≡LC/Ls)随着源的功率的增大而增加，同步辐射峰值光度)随着同步辐射峰值频率的增大而减少.Ghisellini等人拟合了51个耀变体的SED，其结果表明高功率耀变体辐射的能量密度较大[9]，且同步峰值频率辐射的电子能量γp和在Thomson散射截面下磁场和辐射场的能量密度存在负相关性[9-13].导致上述结论的原因是：有效的逆康普顿散射导致较低的电子能量和较大的CD，而低能电子发出较低频率的同步辐射光子.通常把vs-Ls和之间的关系叫做耀变体序列.Abdo等人发现伽马射线光度(Lγ)与伽马射线光子谱指数(Γγ)之间具有相关性[14-15].由于光子谱指数与峰值频率之间有相关性,而伽马射线光度可以表示峰值频率大小[15]，因此，Lγ-Γγ关系可用于耀变体演化序列研究.近几年，关于耀变体序列有不同的观点[16-20]，研究者发现了一些射电光度低且同步辐射峰值频率也低的耀变体[21].同时，高同步辐射峰值频率的FSRQs也被找到[18].还有一些研究者认为，耀变体能谱系列可能是多普勒效应造成的，即认为低同步辐射峰频、高同步辐射光度耀变体的多普勒增亮效应更显著.而Ghisellini等的研究认为：(1)低vs、低Ls的耀变体可能偏离观测者视线方向，弱的聚束效应导致了耀变体的频率和光度都较低；(2)低vs、低Ls的耀变体可能对应质量较小的黑洞，因此喷流将在宽线区内耗散能量，电子在低频率部分辐射且冷却效率高[22]. 本文收集了费米选耀变体的中心黑洞质量和准同时性多波段数据，用对数抛物线拟合其SED并研究耀变体序列.文中取哈勃常数H0=70 km·s-1·Mpc-1，物质能量密度ΩM=0.3，无量纲宇宙学常数ΩΛ=0.7.耀变体是一类快速光变的源，因此，使用同时性多波段数据来拟合SED并与理论模型作比对非常重要.由于观测条件限制，只能选择最大样本费米选耀变体准同时性多波段数据.利用ASDC SED Builder*一种基于线上服务的ASI Science Data Center(ASDC)[21]，http://tools.asdc.asi.it/SED/，收集了129个费米耀变体(其中74个FSRQ，55个BL Lacs)在Planck、WISE、Swift和Fermi(2009年8月-2010年6月)等波段的多波段准同时性观测数据.样本不包括Chen的样本[24]中48个费米LAT亮源AGN(LBAS)和4个费米选窄线Seyfert 1(NLS1)源.逆康普顿散射与同步辐射峰值频率由二次多项式y=ax2+bx+c拟合得到，其中y=logvFv，x=logv.在源坐标系下，其峰值光度和频率由以下两个公式计算得到，和vs,C=(1+z)vp_obss,C，dL是光度距离，z是红移.从文献[25-33]中收集了样本源的黑洞质量，如果一个样本源有多个不同的黑洞质量，则对其取平均值.伽马射线光度和伽马射线光子谱指数来源于文献[34].相应的数据详见文献列表*由于样本比较大所以样本数据和SED图在文中没有给出，如有需要请与作者联系.，其中各栏信息如下：(1)名称；(2)和(3)为同步辐射峰值频率和光度；(4)和(5)为逆康普顿散射峰值频率和光度；(6)红移；(7)伽马射线光子谱指数；(8)伽马射线光度；(9)光学类型；(10)和(11)为黑洞质量及参考文献.3.1 拟合结果有效性测试Ackermann等人最近发布了AGN的费米三期数据(3LAC)，采用了一种基于三次多项式拟合非同时性SED来估算同步辐射峰值频率的方法[35].为了测试所拟合的峰频的有效性，从文献[35]中得到92个blazars(40个FSRQs和52 BL Lacs)的同步辐射峰值频率，并把我们的拟合结果与文献[35]中的同步辐射峰值频率作相关分析，结果在图1中给出.左图为FSRQs，右图为BL Lacs，横坐标为文献[35]中同步辐射峰值频率，纵坐标为我们拟合得到的同步辐射峰值频率，图中黑色虚线为y=x，黑色实线为数据最佳拟合结果，斜率分别为0.874和1.03，皮尔逊概率分别为p=4.91×10-9和p<10-20.从图中可以看到，估算得到的FSRQs的峰值频率多数小于文献[35]中的峰值频率，而BL Lacs的峰值频率多数大于文献[35]中的峰值频率.Ackermann等人认为，使用非同时性多波段数据可能由热辐射/吸积盘辐射导致高估FSRQs的同步辐射峰值频率，而寄主星系的贡献可能低估了对BL Lacs的同步辐射峰值频率.因此，我们用准同时性多波段观测数据拟合得到的vpeak是可靠的.3.2 耀变体序列同步峰值频率vs和同步峰值光度Ls间的相关性如图2所示，图中横坐标为logvs，纵坐标logLs.从图中可以看出随着vs的增大Ls减小，blazar演化遵循从FSRQs到LBLs到HBLs的顺序，沿着这个序列，耀变体的峰值光度在减小，峰值频率在增大.上述结论和利用伽马光子谱指数与伽马射线光度得到的演化序列一致[14-15，38].另外，图2中还呈现出一些低vs、低Ls的耀变体.Ghisellini等人认为，视向偏离或中心有小质量黑洞可能导致低vs低Ls耀变体[22].由于康普顿成分主导度CD是一个和红移无关的量，因此其在演化序列研究中是一个重要的参数.Finke等研究了CD与同步辐射峰值频率间的关系，得出康普顿主导度CD随着同步辐射峰值频率的增加而减少[36].我们研究了CD与同步峰值光度Ls间的关系，结果如图3所示.从图3中可以看出CD随着Ls的增大而增大，这与Fossati等人得到的逆康普顿成分的主导度CD随着源功率的增加而增加[8]的结论一致.如果这些耀变体有相对论性大视角，则会偏向较低频率和较低光度.逆康普顿散射峰值频率和同步辐射峰值频率都依赖聚束效应，因此，rCs=vC/vs和CD与视角无关.由于L∝δ4,因此预料rCs和CD将不依赖.将rCs和CD与的相关性分别呈现在图4、图5中.从图中可以看到大的离散，这个结果不支持低vs、低Ls的耀变体是视向偏离.Ghisellini等人认为低vs低Ls耀变体可能是有小质量黑洞，喷流将在宽线区内耗散能量造成电子的有效冷却和引起低频率和低光度.小质量黑洞也引起康普顿主导度偏低[22].为了调查是否是黑洞造成这些低vs低Ls耀变体，我们从样本中找出黑洞质量.图2中给出的最佳拟合是Ls∝vs-0.39±0.046，因此，参数Lsvs0.39和黑洞质量的相关性可以被用来检测低vs、低Ls的耀变体是否有小质量黑洞，相关结果可以从图6中看到，有明显的正相关，尽管有小的离散，但我们的结果支持低vs、低Ls的耀变体有小质量黑洞.因此，证实了Ghisellini等人认为中心有小质量黑洞可能导致低vs低Ls耀变体[20]的观点.在本文中，把BL Lacs天体分为HBLs和LBLs,从图2耀变体的分布中可以看到大部分LBLs落在FSRQs中，除了LBLs没有或只有弱的发射线这一点外，LBLs和FSRQs具有共同特性，Blazar三个子类在它们的谱特性上显示了一致的连续性，形成了一个序列.在讨论费米选耀变体时，人们有时会使用低同步峰值频率耀变体(LSP)、中同步峰值频率耀变体(ISP)和高同步峰值频率耀变体(HSP)代替FSRQs和BL Lacs[13].研究结果与Padoani等人和Chen的结果一致，尽管Padoani等人的研究基于射电波段到X射线波段，而我们样本数据基于射电到γ射线波段，两个结果都呈现出低vs、低Ls的耀变体.由于缺乏γ射线波段数据所以Padoani等人并没有研究康普顿主导度(CD).Ghisellini等人研究了费米选耀变体并且给出了Lγ和Γγ之间的相关性，认为降低γ射线流量阈值将发现陡谱指数耀变体和低同步辐射光度耀变体[38]，Lγ和Γγ之间有明显的正相关[14].图2中看到logv-log(vLs)之间存在负相关之外还呈现一些低vs、低Ls的耀变体，因此当讨论光子谱指数与峰值频率相关性、伽马射线光度与峰值光度之间的相关性时需要多注意.本文收集了56个耀变体的伽马射线光度和伽马射线光子谱指数，Lγ和Γγ的相关性在图7中给出，从图中可以清晰地看到两者明显的正相关，与Ghisellini等人的结果一致.结果表明是由黑洞质量导致低vs和低Ls而不是聚束效应.需要说明的是，图4、图5中可以看到数据点比较分散，这意味着聚束效应也有一定的作用.很多射电星系已被费米认证，在射电噪类星体统一模型下，blazars天体属于射电星系，但射电星系视角大.Abdo等人在研究中给出了伽马射线光子谱指数与伽马射线光度间的相关(包括射电星系)，可以看出射电星系有较低的光度和光谱较平滑，这与视向偏离的耀变体有低vs和低Ls的假设是一致的[15].同步辐射峰值光度Ls和黑洞质量的相关性在图8中给出，从图中可以看到存在正相关，说明同步辐射峰值光度低的耀变体确实有小质量黑洞.总之，费米选耀变体的logv-log(vLs)之间存在负相关性，Ls随着vs的增大而减小，blazar演化遵循从FSRQs到LBLs到HBLs的顺序.呈现出低vs、低Ls的耀变体的原因可能和小质量黑洞有关而不是弱的聚束效应.【相关文献】[1] 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[22]GHISELLINI G,TAVECCHIO F.The blazar sequence: a new perspective[J].Monthly Notices of the Royal Astronomical Society,2008,387(4):1669-1680.[23]STRATTA G,CAPALBI M,GIOMMI P,et al.The ASDC SED Builder Tool description and tutorial[J].arXiv preprint arXiv:1103.0749,2011.[24]CHEN L,BAI J M.Implications for the blazar sequence and inverse compton models from Fermi bright blazars[J].The Astrophysical Journal,2011,735(2):108.[25]WOO J H,URRY C M,VAN DER MAREL R P,et al.Black hole masses and host galaxy evolution of radio-loud active galactic nuclei[J].The Astrophysical Journal,2005,631(2):762.[26]XIE G Z,ZHOU S B,LIANG E W.The mass-luminosity relation,accretion rate-luminosity relation, and evolutionary sequence of blazars[J].The Astronomical Journal,2004,127(1):53.[27]LIU Y,JIANG D R,GU M F.The jet power, radio loudness,and black hole mass in radio-loud active galactic nuclei[J].The Astrophysical Journal,2006,637(2):669.[28]ZHANG J,LIANG E W,ZHANG S N,et al.Radiation mechanisms and physical properties of GeV-TeV BL Lac objects[J].The Astrophysical Journal,2012,752(2):157.[29]SBARRATO T,GHISELLINI G,MARASCHI L,et al.Th e relation between broad lines and γ-ray lumi-nosities in Fermi blazars[J].Monthly Notices of the Royal Astronomical Society,2012,421(2):1764-1778.[30]CHAI B,CAO X,GU M.What governs the bulk velocity of the jet components in active galactic nuclei?[J].The Astrophysical Journal,2012,759(2):114.[31]LEN-TAVARES J,VALTAOJA E,CHAVUSHYAN V H,et al.The connection between black hole mass and Doppler boosted emission in BL Lacertae type objects[J].Monthly Notices of the Royal Astronomical Society,2011,411(2):1127-1136.[32]SHEN Y,RICHARDS G T,STRAUSS M A,et al.A catalog of quasar properties from Sloan Digital Sky Survey data release 7[J].The Astrophysical Journal SupplementSeries,2011,194(2):45.[33]SHAW M S,ROMANI R W,COTTER G,et al.Spectroscopy of broad-line blazars from1LAC[J].The Astrophysical Journal,2012,748(1):49.[34]ACKERMANN M,AJELLO M,ALLAFORT A,et al.The second catalog of active galactic nuclei dete-cted by the fermi large area telescope[J].The AstrophysicalJournal,2011,743(2):171.[35]ACKERMANN M,AJELLO M,ATWOOD W,et al.The Third Catalog of Active Galactic Nuclei Detected by the Fermi Large Area Telescope[J].arXiv preprintarXiv:1501.06054,2015.[36]FINKE J pton Dominance and the Blazar Sequence[J].The Astrophysical Journal,2013,763(2):134.[37]TAVECCHIO F,MARASCHI L,GHISELLINI G.Constraints on the physical parameters of TeV blazars[J].The Astrophysical Journal,1998,509(2):608.[38]GHISELLINI G,MARASCHI L,TAVECCHIO F.The Fermi blazars' divide[J].Monthly Notices of the Royal Astronomical Society,2009,396(1):L105-L109.。

1汉英力学名词(1993)

BZ反应||Belousov-Zhabotinski reaction, BZ reactionFPU问题||Fermi-Pasta-Ulam problem, FPU problemKBM方法||KBM method, Krylov-Bogoliubov-Mitropolskii method KS[动态]熵||Kolmogorov-Sinai entropy, KS entropyKdV 方程||KdV equationU形管||U-tubeWKB方法||WKB method, Wentzel-Kramers-Brillouin method[彻]体力||body force[单]元||element[第二类]拉格朗日方程||Lagrange equation [of the second kind] [叠栅]云纹||moiré fringe; 物理学称“叠栅条纹”。

[叠栅]云纹法||moiré method[抗]剪切角||angle of shear resistance[可]变形体||deformable body[钱]币状裂纹||penny-shape crack[映]象||image[圆]筒||cylinder[圆]柱壳||cylindrical shell[转]轴||shaft[转动]瞬心||instantaneous center [of rotation][转动]瞬轴||instantaneous axis [of rotation][状]态变量||state variable[状]态空间||state space[自]适应网格||[self-]adaptive meshＣ0连续问题||C0-continuous problemＣ1连续问题||C1-continuous problemＣＦＬ条件||Courant-Friedrichs-Lewy condition, CFL condition ＨＲＲ场||Hutchinson-Rice-Rosengren fieldＪ积分||J-integralＪ阻力曲线||J-resistance curveＫＡＭ定理||Kolgomorov-Arnol'd-Moser theorem, KAM theoremＫＡＭ环面||KAM torusｈ收敛||h-convergenceｐ收敛||p-convergenceπ定理||Buckingham theorem, pi theorem阿尔曼西应变||Almansis strain阿尔文波||Alfven wave阿基米德原理||Archimedes principle阿诺德舌[头]||Arnol'd tongue阿佩尔方程||Appel equation阿特伍德机||Atwood machine埃克曼边界层||Ekman boundary layer埃克曼流||Ekman flow埃克曼数||Ekman number埃克特数||Eckert number埃农吸引子||Henon attractor艾里应力函数||Airy stress function鞍点||saddle [point]鞍结分岔||saddle-node bifurcation安定[性]理论||shake-down theory安全寿命||safe life安全系数||safety factor安全裕度||safety margin暗条纹||dark fringe奥尔-索末菲方程||Orr-Sommerfeld equation奥辛流||Oseen flow奥伊洛特模型||Oldroyd model八面体剪应变||octohedral shear strain八面体剪应力||octohedral shear stress八面体剪应力理论||octohedral shear stress theory巴塞特力||Basset force白光散斑法||white-light speckle method摆||pendulum摆振||shimmy板||plate板块法||panel method板元||plate element半导体应变计||semiconductor strain gage半峰宽度||half-peak width半解析法||semi-analytical method半逆解法||semi-inverse method半频进动||half frequency precession半向同性张量||hemitropic tensor半隐格式||semi-implicit scheme薄壁杆||thin-walled bar薄壁梁||thin-walled beam薄壁筒||thin-walled cylinder薄膜比拟||membrane analogy薄翼理论||thin-airfoil theory保单调差分格式||monotonicity preserving difference scheme 保守力||conservative force保守系||conservative system爆发||blow up爆高||height of burst爆轰||detonation; 又称“爆震”。

吸积模型应用到活动星系核所遇到的困难-星际物理与化学

பைடு நூலகம்

Additional heating to the corona is often assumed. But the mechanism is unclear!
• Temperarure inpedendent on M Ti~1011-1012K，近维里温度 Te~109K • Low L: Eddington ratio≾0.01
Narayan et al. 1998
Comparison of the Spectra
log fν soft
hard log hν
Accretion Energy Radiation Low T, high L, geometrically thin disk ≫ 1: multi-color blackbody Soft state spectra
•
Advection-Dominated Accretion Flow (ADAF)
Spectrum from Standard Disk
Multi-colour Blackbody
spectra (optically thick)

kT ~1 keV ( BH binaries)
Log n f(n)
kT ~ 10 eV (AGN)
Log n
Spectrum from ADAF
活动星系核（AGN)
• 活动星系核：星系的核心有剧烈电磁辐射 • 提供活动星系核巨大能源的发动机—黑洞吸积过程星系 • 低吸积率的黑洞：正常星系 • 高吸积率的黑洞：活动星系核
活动星系核类星体黑洞
吸积的基本模式
•
The standard accretion disk (Shakura & Sunyaev 1973)

i Running Title Accretion in the Kuiper Belt

Accretion in the Edgeworth-Kuiper Belt: Forming 100{1000 km Radius Bodies at 30 AU and Beyond
S. Alan Stern Geophysical, Astrophysical, & Planetary Sciences Space Science Department Southwest Research Institute Boulder, CO 80302 USA and Joshua E. Colwell Laboratory for Atmospheric and Space Physics University of Colorado Boulder, CO 80309 USA
S. Alan Stern Space Science Department Southwest Research Institute 1050 Walnut Street, Suite 426 Boulder, CO 80302 303] 546-9670 (voice) 303] 546-9687 (fax) astemploy a time-dependent collisional evolution code to study the conditions under which the 50{200 km radius Edgeworth-Kuiper Objects (EKOs) in the region between 30 and 50 AU (now called the Edgeworth-Kuiper Belt, or EKB; Edgeworth 1943, 1949; Kuiper 1951) were formed. Assuming that these bodies were created by pairwise accretion from 1 to 10 km building blocks, we nd that three conditions were required, namely: (i) at least 10 M and more likely 35 M of solids in the primordial 30 to 50 AU zone, (ii) mean random orbital eccentricities of order 0.002 or smaller, and (iii) mechanically strong building blocks. Furthermore, we nd that the accretion of 100{200 km radius bodies in the 30 to 50 AU region from collisions among a starting population of 1 to 10 km building blocks required 108 {109 years, with the lower range only being reached for 30 to 50 AU zone masses approaching 100 M of solids or mean random orbital eccentricities <0.005 (which we do not believe is realistic after gas dissipation). Therefore, unless accretion had already produced many building blocks signi cantly larger than 10 km in diameter at the time the nebular gas was removed, our results also indicate that Neptune did not form on a timescale much shorter than 70 Myr, and could well have required many hundreds of Myr to approach its nal mass. We also explore the growth of Pluto-scale (i.e., radius 1000{ 1200 km) objects in the 30 to 50 AU region under a variety of assumptions. We further nd that once 300 hundred kilometer radius objects were formed, the growth of 1000 km radius and larger objects occurs relatively easily and comparatively quickly. The lack of many Plutos in the 30 to 50 AU zone therefore argues strongly that growth was terminated in that region rather abruptly at the time the presently observed population of 100{200 km radius EKOs were being completed. In the region beyond 50 AU where Neptune's dynamical in uence was much reduced, model runs yield 100 km to 1000 km radius, and perhaps even larger bodies which should be detectable with on-going or soon-to-be started surveys. We suggest that if dynamical conditions did not remain calm enough to allow Pluto itself to be grown in the 30 to 50 AU zone before perturbations from Neptune created a dynamically erosive, low-mass environment there, then it may be that Pluto was grown beyond the in uence of Neptune's perturbations and later transported inward, perhaps in part via the Charon-forming collision.

Quark-nova remnants IIIApplication to RRATs

a r X i v:0802.3929v 1 [a s t r o -p h ] 26 F e b 2008Astronomy &Astrophysics manuscript no.(will be inserted by hand later)Quark-nova remnants III:Application to RRATsRachid Ouyed 1,Denis Leahy 1,Brian Niebergal 1,and Youling Yue 1,21Department of Physics and Astronomy,University of Calgary,2500University Drive NW,Calgary,Alberta,T2N 1N4Canada 2Astronomy Department,School of Physics,Peking University,Beijing 100871,Chinarecieved /acceptedAbstract.This is the third paper of a series of papers where we explore the evolution of iron-rich ejecta from quark-novae.Inthe ﬁrst paper,we explored the case where a quark-nova ejecta forms a degenerate shell,supported by the star’s magnetic ﬁeld,with applications to SGRs.In the second paper we considered the case where the ejecta would have su ﬃcient angular momen-tum to form a degenerate Keplerian torus and applied such a system to two AXPs,namely 1E2259+586and 4U0142+615.Here we explore the late evolution of the degenerate torus and ﬁnd that it can remain unchanged for ∼106years before it becomes non-degenerate.This transition from a degenerate torus (accretion dominated)to a non-degenerate disk (no accretion),occurs about 106years following the quark-nova,and exhibits features that are reminiscent of observed properties of ing this model,we can account for the duration of both the radio bursts and the quiet phase,as well as the observed radio ﬂux from RRATs.The unique on and o ﬀactivity of the radio pulsar PSR B1931+24is similar to that of “old RRATs”in our model.For old RRATs,in our model,the spin-down rate during the radio loud phase is about a factor 1.6larger than when it is quiet,remarkably similar to what has been measured for PSR B1931+24.We discuss a connection between XDINs and RRATs and argue that some XDINs may be “dead RRATs”that have already consumed their non-degenerate disk.Key words.accretion,accretion disks –(stars:)pulsars:general –dense matter –X-rays:bursts –Elementary particles1.IntroductionMcLaughlin et al.(2006)have recently reported the detection of eleven “Rotation RAdio Transients”,or “RRATs”,charac-terised by repeated,irregular radio bursts,with burst dura-tions of 2-30ms,and intervals between bursts of ∼4min to ∼3hr.The RRATs are concentrated at low Galactic lat-itudes,with distances implied by their dispersion measures of ∼2−7kpc.For ten of the eleven RRATs discovered by McLaughlin et al.(2006),an analysis of the spacings between repeat bursts reveals an underlying spin period,P ,and also inthree cases,a spin period derivative,˙P.The observed periods fall in the range 0.4s <P <7s,which generally overlap with those seen for the radio pulsar population.Since August 2003,all the sources have been reobserved at least nine times at intervals of between one and six months (i.e.they show spo-radic radio bursting for years).All have shown multiple bursts,with between 4and 229events detected in total from each ob-ject (see Table 1in McLaughlin et al.(2006)).For the threeRRATs with values measured for both P and ˙P,a characteristic age,τc ,and a dipole surface magnetic ﬁeld,B ,can be inferred and are listed in Table 1.The total number of these objects is a few times that previously estimated for all radio pulsars (McLaughlin et al.(2006)).1No radio emission of any kind has been reported from XDINs;recent observations (Bradley 2006)show no RRAT-like radio bursts toward RX J0720.4–3125(or toward magnetars).2Ouyed et al.:Quark-nova remnants IIInous(L X∼1031−1032ergs s−1)than J1819–1458.However, the measured period derivatives of two XDINs(RBS1223and RX J0720.4–3125;Kaplan&van Kerkwijk2005a,2005b), and the detection of possible proton cyclotron lines in their spectra(van Kerkwijk2004),imply magneticﬁeld strengths similar to those of J1819–1458.The immediate questions are (i)why only one RRAT shows X-rays as the XDINs do, and(ii)why XDINs do not show any radio emission(latest searches for pulsed and bursty radio emission from XDINs led to no detection despite of their proximity compared to RRATs; Kondratiev et al.(2007)).RRAT J1819–1458properties in the X-ray show similari-ties to those of radio pulsars with ages around100kyr.For ex-ample,PSR J0538+2817is30kyr old and has kT∞=160eV, while PSR B0656+14is110kyr old and has kT∞=70eV (see Reynolds et al.(2006)for details).However,the inferred surface magneticﬁeld strength of RRATs is at least an order of magnitude greater than the radio sources.Two radio pul-sars with comparable magneticﬁelds that have been detected in X-rays are PSRs J1718–3718(Kaspi and McLaughlin(2005)) and J1119–6127(Gonzalez et al.(2005)).These sources show temperatures(kT∼150−200eV)and luminosities(∼1032−1033ergs s−1)comparable to that of RRAT J1819−1458, although both sources are probably much younger(35and 1.7kyr,respectively)and,in contrast to RRAT J1819–1456, have X-ray luminosities less than their spin-down luminosities.1.1.Literature explanationsThe models proposed so far in the literature can be classi-ﬁed into the following categories:(i)Extreme pulses from distant pulsars(Weltevrede et al.(2006)),similar to the pulses seen from the nearby pulsar B0656+14.(ii)Re-activated ra-dio pulsars(Zhang et al.(2007))where the RRATs are pul-sars that are no longer active,but for which a temporary“star spot”with multipoleﬁeld components emerges above the sur-face.This magneticﬁeld component could temporarily reac-tivate the radio beaming mechanism,producing the observed bursts.(iii)Nulling pulsars viewed from the opposite direc-tion(Zhang et al.(2007)).In this model,RRATs are normal pulsars with their magnetic poles not aligned favorably for de-tection,but undergo an magnetic reversal and so occasionally produce emission that can be observed.(iv)Sporadic accre-tion(Cordes and Shannon(2006),Li(2006)).Here the RRAT mechanism might be produced by interaction of the neutron star with an equatorial fallback disk or with orbiting circum-polar debris.Accretion from a disk should quench the radio emission mechanism,but sporadic drops in the accretion rate could allow the radio beam to turn on for a fraction of a sec-ond,producing the RRAT phenomenon(Li(2006)).In this paper,the third of a series,we present an alternative scenario that incorporates SGRs,AXPs,XDINs,and RRATs into one family of compact objects.In our model,SGRs,AXPs, XDINs,and RRATs are all strange quark stars that diﬀer,for the most part,only by age.This paper is organized as follows:Sect.2describes our strange quark star model.Sect.3discusses the implications of the evolution of debris material left over from the birth of the strange quark star.In Sect.4we show how after roughly one million years the debris can become responsible for RRAT be-haviour.Then,in Sect.5we describe observations and make predictions using our model.We discuss some further implica-tions of our model in Sect.5,and lastly we conclude in Sect.6.2.Our modelIn the quark-nova(QN)picture(Ouyed et al.(2002), Ker¨a nen et al.(2005);hereafter ODD and KOJ respectively) the core of a neutron star,that undergoes the phase transition to the quark phase,shrinks in a spherically symmetric fashion to a stable,more compact strange matter conﬁguration faster than the overlaying material(the neutron-rich hadronic envelope) can respond,leading to an eﬀective core collapse.The core of the neutron star is a few kilometers in radius initially, but shrinks to1-2km in a collapse time of about0.1ms (Lugones et al.(1994)).The gravitational potential energy released(plus latent heat of phase transition)during this event is converted partly into internal energy and partly into an outward propagating shock wave that imparts kinetic energy to the overlying material.As described in a series of papers(Ouyed et al.(2007a), Ouyed et al.(2007b);hereafter referred to as OLNI and OLNII),during a quark-nova the degenerate crust of a neu-tron star is blown oﬀ,leaving behind a strange quark star(QS) surrounded by left over,highly-metallic degenerate matter.In OLNI we discussed the case where the ejected crust had insuﬃ-cient angular momentum to escape the QS’s gravitational pull, and so would either balance with the QS’s magneticﬁeld and form a co-rotating shell,or fall back entirely.Then in OLNII the case where the QS was born with a suﬃcient spin-period to impart the ejected crust with enough angular momentum to form a degenerate torus.In this paper,we explore the result of this degenerate torus after enough time has passed for it to expand to densities where it is non-degenerate.2.1.Emission:Vortex Expulsion vs.AccretionIn our model the QS is in the ground Color-Flavor Locked (CFL)phase and so behaves as a type II superconductor, wherein a rotationally-induced lattice forms inside the star.As the star spins down the magneticﬁeld,which is conﬁned to ex-ist only within the vortices,is also expelled and the subsequent magnetic reconnection leads to the production of X-rays.The luminosity from vortex expulsion is given in OLNI to be,L X,v≃2.01×1034erg s−1ηX,0.1˙P2−11.(1) Here the subscript v stands for“vortex”andηX,0.1is the ef-ﬁciency parameter,in units of0.1,inherent in the conversion from magnetic energy to observed radiation.The spin-down rate,˙P is given in units of10−11s s−1.In this paper the degenerate material ejected during the quark-nova is imparted with suﬃcient angular momentum to form a degenerate torus,in which case accretion from this torus can result in emission that outshines the emission due to vortexOuyed et al.:Quark-nova remnants III3 expulsion.The condition on the initial spin-period of the QS atbirth is then an upper limit(from OLNII),P0<P0,max=2.5ms B3/20,15R9/2QS,10µ63.3M4QS,1.4.(3)Hereη0.1is the eﬃciency of conversion of accreted material into X-ray emission in units of0.1,R t,15is the radial distance of the torus in units of15km,µ3.3is the mean molecular weight per electron in units of3.3(the quiescent phase value),and M QS,1.4is the mass of the QS in units of1.4solar masses.2.2.Degenerate Torus propertiesThe radius of the torus can be found from the initial spin-period of the QS as was done in OLNII,R t=15kmα−8/30.3P2/30,msM1/3QS,1.4,(4)where the initial period is in units of milliseconds.By assuming a constant accretion rate,the evolution of the torus mass can be determined,and the time needed for the torus to reach a density where it becomes non-degenerate is,τt≃m0η30.1R6t,15,(5)with˙m t being the torus accretion rate given by eq(10)in OLNII. The mass and density of the torus over time are plotted in Fig. 1.While the torus is above degenerate densities,diﬀusion by the QS’s magneticﬁeld into the inner walls of the torus(lead-ing to bursting accretion events;see OLNII)increases the inner radius.We can estimate the change inner radius of the torus, R in=R t+∆R t(see Appendix A),by using the typical time needed for magneticﬁeld diﬀusion into the torus,τB,and the typical depth into the torus at which shear stresses become great enough such that accretion can proceed,∆r w.Both of these parameters were estimated in OLNII(eqns.17&5).We ﬁnd that the change in inner radius over the span of the lifetime of the torus is negligible.While the magneticﬁeld slowly consumes the torus’inner edge,its outer edge moves outward(due to viscosity from par-ticle collisions within a degenerate ideal gas)at a rate given by equation(A.7)in OLNII.After the time where the torus ex-pands to densities at which it becomes non-degenerate,τt,the torus will have extended to an outer radius of,R out≃207km m1/20,−7M2QS,1.4µ3q,3.3τv14Ouyed et al.:Quark-nova remnants IIIand ˙PQS (t )=P QS ,0τv −23κP ˙P PB 2=P 0B 20.(10)where κ=8.8×1033G 2s −1.The former equation describesmagnetic ﬁeld decay while the latter allows us to link the initial and current conditions.By making use of the above equations and equation (4),the star’s period,period derivative,and magnetic ﬁeld can be deter-mined at the time,τt ,when the transition to a non-degenerate torus is made (with τt >>τv ),P trans .≃2.3sµ2q ,3.3B 4/30,15M 1/2QS ,1.4R 10/3QS ,10µ4q ,3.3B 2/30,15M 2QS ,1.4R 32/3QS ,10(11)Combining equation above with eq.(10)we ﬁndB 0,15≃2.1f ∗P 35/38trans .,3˙P 13/38trans .,−13with P trans .and ˙P trans .in units of 3s and 10−13s s −1,respectively,whilef ∗=(R 10QS ,10µq ,3.3M −1/3QS ,1.4η−6/130.1α8/30.3)−39/76is of order unity.Equations (4)and (11)give us R t ,15and P 0,ms .With P 0,ms ,B 0,15and eq(2)together,we solve for m 0,−7.Table 2shows birth parameters of the 3RRATs studied here for α=0.3.We ﬁnd an initial magnetic ﬁeld 1014<B 0G <1016and sub-millisecond birth periods.Also listed in Table 2are the torus’initial inner radius,R t ,and initial mass,m 0which implies 10−8M ⊙<m 0<10−5M ⊙and 10<R t (km)<35.3.2.Non-degenerate disk propertiesThe non-degenerate disk’s thickness at a radial distance r is given by H d =v 2th ./g ,where g =GM QS H d /r 3is the e ﬀective gravity at r ,and v th =µ1/2q ,3.3M 1/2QS ,1.4,(12)while the disk’s thickness is a few hundred meters at its outeredge,R out .However,as we show below,the disk quickly cools reducing its temperature and thus its thickness until it solidiﬁes.The density would continue dropping below ρnd until it reaches the density of normal iron ρFe ∼7.86g cm −3.The thermal evolution of the disk is given by,C v∂T4πL X ,v −L BB ,t ,(13)where Ωd =(2πH d /R d )is the solid angle extended by the diskwith thickness H d .In the above,C v is the torus heat capacitywhile L BB ,t ≃2πR 2out σT 4is the torus blackbody cooling.The disk cools,while H d decreases,yielding smaller Ωd thus further cooling.However,the decrease in H d stops when the disk gas condenses which occurs at a temperature of T Fe ,cond .≃0.265eV (CRC tables),since the disk composition is dominated by iron (see OLNI and OLNII).To ﬁnd the disk’s thickness at the condensation point we impose conservation of surface density so that H d ρnd =H d ,cond .ρFe ;we assume the surface density constant during the gas phase since the gas phase is short-lived so that H d ,cond .∼0.12cmR 3/2in ,100R 1/2out ,1000µ1/8q ,3.3M 1/8QS ,1.4,(15)where R out is given in units of 1000km.For the 3RRATs in the order listed listed in Tables we get,∼6meV ,∼86meV and ∼5.7meV ,respectively.At temperatures below 0.1eV ,con-densed iron is in the form of ferrite,or α-iron,a body-centered cubic structure.4.RRAT-like sporadic behaviorThe inner disk will be slowly penetrated by the QS’s magnetic ﬁeld on timescales determined by the induction equation,∂B4πσ∇2B .(16)Here,σ=n e ,th e 2λe /(m e v rms )and λe =1/(n e ,th σT ),where n e ,th is the number density of thermal electrons in the disk,σT is the Thompson scattering cross-section,and the root-mean-square electron velocity is v rms =v th .Therefore,the time needed for the magnetic ﬁeld to penetrate to a depth of the order of H d ,cond .,τB ∼1.3hrsR 3in ,100v ﬀ∼2.3msR 3/2in ,1002GM QS /R in is the free-fall velocity.Ouyed et al.:Quark-nova remnants III 5Table 1.Observational properties of the RRATs.RRAT J1317–5759 2.60.126<7.5×1032100.221.4×1031RRAT J1819–1458 4.3 5.763.3×103330.0575.6×1031RRAT J1913+33330.920.0787<9.4×103420.212.5×1031Star TorusSource B 0(1015G)P 0(ms)˙P0(10−6s s −1)R t (km)m 0(10−7M ⊙)Table 3.RRATs era parameters predicted by our model (for α=0.3).SourceL X ,v (erg s −1)T BB (eV)R in (km)R out (km)T d (meV)t on =t acc .(ms)t o ﬀ=τB (hr)L Radio (erg s −1)L X fromeq(1),T BB from eq(20),R in from eq(A.2),R out from eq(6),T d from eq(15),t acc from eq(18),τB from eq(17),L radio from eq(19)withA casc .=102.Table 4.(Torus and RRATs)Era duration predicted by our model (for α=0.3).RRAT J1317–57592.2×1067.9×105RRAT J1819–14587.9×1043.2×103RRAT J1913+33331.3×1064.1×105µq ,3.3M 1/2QS ,1.4.(19)As shown in Table 3,the observed radio ﬂuxes can be ac-counted for consistently for all 3candidates if an ampliﬁcationfactor of the order of 102is assumed (a multiplicity number for secondary pairs can be as high as 105;e.g.Melrose 1995and references therein).We expect the medium to be transparent to radio since it would most likely be emitted much above the polar region.We assume all radio bursts originate at the same longitude on the quark star.This guarantees that the bursts are modulated at the spin period,consistent with the assumptionrequired to derive P and ˙Pfrom the observations of RRATs.Table 3lists parameters derived from our model as com-pared to observations which shows some encouraging agree-ment with measured radio burst timescale and quiet phases.The case of RRAT J1819−1458shows smaller bursting time and quiet time,as well as smaller radio ﬂuxes probably because our model somehow predicts a smaller R in than the actual one.For example,using R in ∼80km instead of 33km (i.e.α∼0.2instead of α=0.3)we ﬁnd much better agreement.5.Model Predictions5.1.The two blackbodiesOne of the key aspects of the transition from degenerate torus toa non-degenerate disk is the shutting-o ﬀof accretion thus elim-inating one of the two BBs inherent to the accretion era (see OLNII).RRATs should show one dominant blackbody emis-sion at a temperature given byT BB =(L X ,v /(4πR 2QS σ))1/4∼0.04keV ˙P 1/2−13,(20)where ˙Pis in units of 10−13s s −1,and a dimmer one at T d rep-resenting the remnant iron disk.In particular,for RRAT J1819-1458for which ˙P−13=5.76we predict T BB ∼0.1keV and6Ouyed et al.:Quark-nova remnants III another blackbody in infrared at∼940K.The blackbody tem-perature for the other two candidates are listed in Table3whichshows a BB temperature in the ultra-violet region and disk tem-peratures in the Infrared,50-100K.5.2.Absorption linesSince the disk is cooler than the star(T d<T BB),it will act asan absorber of the blackbody emission.The column density inthe inner part of the disk isN d,Fe=H d,cond.ρFeM QS,1.4R2QS,10.(22)Compared to˙P trans.we get∆˙PP0,ms R13/3t,15(23)As can be seen from the equation above the increase in spin-down rate when radio is on is negligible is negligible.However,as we show below,as the RRAT ages the increase in spin-downrate during the on period is noticeable.5.4.Old RRATs:the case of PSR B1931+24A simple estimate of the lifetime of a RRAT can be found byintegrating equation dt/dr≃τB/H d,cond.,from R in to R out,togetτrrat∼1.6×106yrs R5/2out,1000M QS,1.4R2QS,10.(26)The spin-down rate during the on era is thus˙P on=˙P oﬀ+∆˙P.The observed periodicity of the radio-onand radio-oﬀrecurrence of PSR B1931+24(Kramer et al.(2006)Kramer,Lyne,O’Brien,Jordan,and Lorimer)is diﬃcult to explain in any scenarioconsidering an isolated pulsar,and asKramer et al.(2006)Kramer,Lyne,O’Brien,Jordan,and Lorimerpointed out,the short shut-oﬀtime of less than10seconds,rules out possible scenarios like precession of the neutron star.Furthermore,the facts that the oﬀperiods of PSR B1931+24areﬁve orders of magnitude longer than typical nullingperiods,that the activity pattern is quasi-periodic and that nota single null has been observed during on periods stronglyargues against the nulling scenario.It appears however thatmagnetospheric conditions are suﬃcient to explain the changein the neutron star torque,but it is not clear what determinesthe∼30-40days periodicity or what could be responsible forchanging the plasmaﬂow in the magnetosphere,in particularin this quasi-periodic fashion.Assuming PSR B1931+24(P=0.813s)is an old RRATthat is currently consuming the outermost part of its disk,ourmodel implies a corresponding spin-down frequency during theoﬀperiod of˙νoﬀ=−ν2˙P oﬀ≃−2.6×10−14Hz s−1,whileduring the on period the increase in spin-down frequency is∆˙ν≃−1.7×10−14Hz s−1.It means an increase of about65%inspin-down rate during the on period so that˙νon≃−4.3×10−14Hz s−1.Our numbers are remarkably similar to what has beenmeasured for PSR B1931+24.For this source the observedspin-down rate the pulsar rotation slows down50%faster whenit is on than when it is oﬀwith˙νoﬀ=−1.08×10−14Hz s−1while˙νon=−1.63×10−14Hz s−1.For PSR B1931+24the quiet period last for∼30-40dayswhile the active period is of the order of∼5days.In ourmodel,the timescales for the bursting and quiet phases in oldRRATs are t acc.∼70ms andτB∼50days.While the oﬀperiodis again remarkably similar to the observed one,the observedon period of∼5days is diﬃcult to reconcile with our modelwhich predicts∼70ms.It is worth noticing however thatKramer et al.(2006)Kramer,Lyne,O’Brien,Jordan,and Lorimerhave been able to observe one switch from on to oﬀthat oc-curred within less than10seconds(although switches betweenthe on and oﬀstates are rare events).This number is rathercloser to t acc.in our model than to the radio-on timescale.Itmight be possible that another mechanism is at play delayingthe accretion onto the pole.Finally,the corresponding radioﬂux in our model is estimated to be L radio∼1034erg s−1.Ouyed et al.:Quark-nova remnants III75.5.XDINs and“dead”RRATs in our modelSome XDINs might be direct descendants of QSs with shells (SGRs in our model;see Appendix B).These would have evolved along the“vortex”band since their birth.However,as illustrated in Figure2,some XDINs might be dead RRATs. That is,they fell into the“vortex”band following torus consumption and eventually evolved along the band.The main diﬀerence between SGR-descendents XDINs and AXP-descendent XDINs(dead RRATs)is the possibility of remnant disk surrounding the AXP-descendent XDINs.We thus expect some of the XDINs to share some common properties with RRATs such as optical(or Infrared)excess.Broad absorption lines,similar to those seen for RRAT J1819−1458,have been observed for six out of seven XDINs(van Kerkwijk&Kaplan 2007;Haberl2007).We have already argued in OLNII that these absorption lines are caused by absorption from an old, cold disk(see also§5.2).Atﬁrst glance the birth rates of RRATs as descendants of AXPs in our model,appears to be too low to explain the high inferred population(McLaughlin et al.2006)of a few times (∼5)that of radio pulsars.Taking the RRAT birth rate to be that of AXPs,i.e.∼1/(500yrs)(Leahy&Ouyed2007and ref-erences therein),implies a birth rate of∼1/10of that of radio pulsars.In order to have a RRAT population∼5times the ra-dio pulsar population,the RRAT lifetime must be∼50times longer than that of a radio pulsar(∼106yrs).In our model,the RRAT lifetime estimate,given by eq.(24),is of order a few mil-lion years.However,this is strongly dependent on R out.An R out larger by a factor of4(caused by diﬀerent eﬀects;see eq.(6)) gives a large enough RRAT lifetime,i.e.a low enough birth rate.6.ConclusionHere we have presented a model where RRATsﬁt in natu-rally as descendants of AXPs with degenerate iron-rich disks, with the transition occurring when the disk becomes non-degenerate.These AXPs in our model are the quark star rem-nants of quark-novae,surrounded by the ejected former neu-tron star’s crust material.The evolution of quark stars as they spin down is summarized in the L X-˙P diagram in Figure2(see appendix B).Although this model is speculative,it can ex-plain many features of SGRs and AXPs(OLNI and OLNII) and RRATs.Acknowledgements.Youling Yue thanks the 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accre-tion events;see OLNII)increasing the inner radius from R t to8Ouyed et al.:Quark-nova remnants IIISGRs (QS+shell)AXPs (QS+torus)XDINs RRATsS t a t i s t i c a l d e t e c t i o n L i n e (S G R s )T r a n s i t i o n t o a c c r e t i o n d o m i na t eder a(A X P s )L X (e r g s -1)P (s s -1).S t a t i s t i c a l d e t e c t i o n L i n e (A X P s )10-1110-1010341031103610-13Fig.2.These ﬁgures depict the important regimes of SGRs,AXPs,RRATs,and XDINs in our model (see app.B),the left panel being an evolutionary illustration and the right panel contains observational data.There are two important transitions to notefor AXPs (torus-harbouring quark stars)in our model,the ﬁrst occurs when the star spins-down to ˙P≈10−11s /s.At this point the dominant luminosity mechanism changes from being magnetic ﬂux (vortex)expulsion to accretion.The second transition,occuring ≈106yrs after the ﬁrst transition,happens when the torus decreases in density enough to become non-degenerate.At this point accretion is no longer e ﬀective,and the star’s emission is again dominated by vortex expulsion.The left-over non-dgenerate disk is slowly consumed resulting in RRAT-like behaviour.Shown in the right panel are the three RRATs studied here,two of which we argue have spent long enough time in the tours (accretion)era before transiting back to having their emission dominated by vortex expulsion.The arrows depicts the large change in their X-ray luminosity during the transition.Also,in the right panel are the 2knowns AXP transients (AXP 1E 1547−54and AXP XTE J1810−197),which fall within the ﬁrst transition,from vortex expulsion to accretion dominated emission.R in =R t +∆R t at τtorus .We can estimate ∆R t by integrating dt /dr =τB /∆r w up to τtorus ,with τB and ∆r w given by eq(5)and (17)in OLNII,respectively.We get95R t ,15 73/24−1,(A.1)which leads to 2cases∆R t ,15≃31R 132/73t ,15If R t ,15≤1.7Appendix B:The L X -˙Pdiagram in our model B.1.Transition from accretion dominated to RRAT eraThe left panel in Figure 2illustrates the torus and RRAT eras in the L X -˙Pdiagram.In the torus phase (lasting roughly 105yrs;eq.5)the X-ray luminosity is dominated by accretion from the torus,during which time the source spins down due to mag-netic braking,thus,evolving horizontally towards smaller ˙P.As the torus is accreted its density decreases (due to accretionand viscous spreading),until it reaches non-degenercy densi-ties,at which point accretion becomes much less e ﬃcient,and the sources luminosity becomes dominated by vortex (mag-netic ﬂux)expulsion from the star.This results in a decrease in luminosity,as is illustrated in ﬁgure 2by the vertical lines.It is this transition,we argue,that is responsible for producing RRAT behaviour.B.2.Transition from vortex to accretion dominated eraBefore the torus phase,there exists a critical ˙Pat which a transi-tion from vortex dominated to accretion dominated luminosity will occur.This critical value is found when L v ≃L acc ,˙Pcrit ∼3×10−12s s−1η20.1R 3t ,15。

第三章恒星的演化

2．较高质量 (M > ~2M⊙) 恒星的演化 (20.4, 21.221.3) (1) 与低质量恒星演化的主要区别恒星内部的H燃烧通过CNO循环进行，内部温度更高，辐射压对维持恒星的力学平衡起更大的作用，主序寿命更短。 He核不再是简并的，C和更重元素的燃烧可以进行。核心区核反应产生的能量主要以对流的方式向外传递。
第三章恒星的演化
§3.1 主序星的演化 §3.2 恒星主序后的演化 §3.3 恒星演化的观测证据 §3.4 密近双星的演化
§3.1 主序星的演化
(20.1)
1. 恒星演化的基本原理
恒星在一生的演化中总是试图处于稳定状态（流体静力学平衡和热平衡）。当恒星无法产生足够多的能量时，它们就无法维持热平衡和流体静力学平衡，于是开始演化。
L ~ M 2.5-4, R ~ M 0.5-1
主序星的演化
(1) 零龄主序 (zero age main-sequence star, ZAMS) 刚刚开始核心H燃烧的恒星，在H-R图上占据主序带的最左侧。 (2) 演化时标 ——核反应 (4 1H→4He +γ) 时标 tn＝η△Mc2/L ≈(1010 yr) (M/M⊙) (L/L⊙)-1 ≈(1010 yr) (M/M⊙)-2.5 for M > M⊙ or (1010 yr) (M/M⊙)-2 for M < M⊙
(4) 特大质量恒星的演化
星风引起的质量损失和恒星演化。高光度恒星通常有很强的星风～10-6-10-4 M⊙yr-1 如沃尔夫-拉叶(WR)星。演化过程 O型星→蓝超巨星→（红超巨星）→WR星→Ib/Ic型超新星 + 中子星/黑洞
小结：
不同质量恒星的演化结局
恒星初始质量 (M⊙) < 0.01 0.01 < M < 0.08 0.08 < M < 0.25 0.25 < M < 8 8 < M < 12 12 < M < 40 M > 40 演化结局行星褐矮星 He白矮星 CO白矮星? ONeMg白矮星? 超新星→中子星? 超新星→黑洞?

天文学专用词汇

Galactic aggregate 银河星集Galactic astronomy 银河系天文Galactic bar 银河系棒galactic bar 星系棒galactic cannibalism 星系吞食galactic content 星系成分galactic merge 星系并合galactic pericentre 近银心点Galactocentric distance 银心距galaxy cluster 星系团Galle ring 伽勒环Galilean transformation 伽利略变换Galileo 〈伽利略〉木星探测器gas-dust complex 气尘复合体Genesis rock 创世岩Gemini Telescope 大型双子望远镜Geoalert, Geophysical Alert Broadcast 地球物理警报广播giant granulation 巨米粒组织giant granule 巨米粒giant radio pulse 巨射电脉冲Ginga 〈星系〉X 射线天文卫星Giotto 〈乔托〉空间探测器glassceramic 微晶玻璃glitch activity 自转突变活动global change 全球变化global sensitivity 全局灵敏度GMC, giant molecular cloud 巨分子云g-mode g 模、重力模gold spot 金斑病GONG, Global Oscillation Network 太阳全球振荡监测网GroupGPS, global positioning system 全球定位系统Granat 〈石榴〉号天文卫星grand design spiral 宏象旋涡星系gravitational astronomy 引力天文gravitational lensing 引力透镜效应gravitational micro-lensing 微引力透镜效应great attractor 巨引源Great Dark Spot 大暗斑Great White Spot 大白斑grism 棱栅GRO, Gamma-Ray Observatory γ射线天文台guidscope 导星镜GW Virginis star 室女GW 型星habitable planet 可居住行星Hakucho 〈天鹅〉X 射线天文卫星Hale Telescope 海尔望远镜halo dwarf 晕族矮星halo globular cluster 晕族球状星团Hanle effect 汉勒效应hard X-ray source 硬X 射线源Hay spot 哈伊斑HEAO, High-Energy Astronomical 〈HEAO〉高能天文台Observatoryheavy-element star 重元素星heiligenschein 灵光Helene 土卫十二helicity 螺度heliocentric radial velocity 日心视向速度heliomagnetosphere 日球磁层helioseismology 日震学helium abundance 氦丰度helium main-sequence 氦主序helium-strong star 强氦线星helium white dwarf 氦白矮星Helix galaxy （NGC 2685 ）螺旋星系Herbig Ae star 赫比格Ae 型星Herbig Be star 赫比格Be 型星Herbig-Haro flow 赫比格-阿罗流Herbig-Haro shock wave 赫比格-阿罗激波hidden magnetic flux 隐磁流high-field pulsar 强磁场脉冲星highly polarized quasar （HPQ ）高偏振类星体high-mass X-ray binary 大质量X 射线双星high-metallicity cluster 高金属度星团;高金属度星系团high-resolution spectrograph 高分辨摄谱仪high-resolution spectroscopy 高分辨分光high - z 大红移Hinotori 〈火鸟〉太阳探测器Hipparcos, High Precision Parallax 〈依巴谷〉卫星Collecting SatelliteHipparcos and Tycho Catalogues 〈依巴谷〉和〈第谷〉星表holographic grating 全息光栅Hooker Telescope 胡克望远镜host galaxy 寄主星系hot R Coronae Borealis star 高温北冕R 型星HST, Hubble Space Telescope 哈勃空间望远镜Hubble age 哈勃年龄Hubble distance 哈勃距离Hubble parameter 哈勃参数Hubble velocity 哈勃速度hump cepheid 驼峰造父变星Hyad 毕团星hybrid-chromosphere star 混合色球星hybrid star 混合大气星hydrogen-deficient star 缺氢星hydrogenous atmosphere 氢型大气hypergiant 特超巨星Ida 艾达（小行星243号）IEH, International Extreme Ultraviolet 〈IEH〉国际极紫外飞行器HitchhikerIERS, International Earth Rotation 国际地球自转服务Serviceimage deconvolution 图象消旋image degradation 星象劣化image dissector 析象管image distoration 星象复原image photon counting system 成象光子计数系统image sharpening 星象增锐image spread 星象扩散度imaging polarimetry 成象偏振测量imaging spectrophotometry 成象分光光度测量immersed echelle 浸渍阶梯光栅impulsive solar flare 脉冲太阳耀斑infralateral arc 外侧晕弧infrared CCD 红外CCDinfrared corona 红外冕infrared helioseismology 红外日震学infrared index 红外infrared observatory 红外天文台infrared spectroscopy 红外分光initial earth 初始地球initial mass distribution 初始质量分布initial planet 初始行星initial star 初始恒星initial sun 初始太阳inner coma 内彗发inner halo cluster 内晕族星团integrability 可积性Integral Sign galaxy （UGC 3697 ）积分号星系integrated diode array （IDA ）集成二极管阵intensified CCD 增强CCDIntercosmos 〈国际宇宙〉天文卫星interline transfer 行间转移intermediate parent body 中间母体intermediate polar 中介偏振星international atomic time 国际原子时International Celestial Reference 国际天球参考系Frame （ICRF ）intraday variation 快速变化intranetwork element 网内元intrinsic dispersion 内廪弥散度ion spot 离子斑IPCS, Image Photon Counting System 图象光子计数器IRIS, Infrared Imager / Spectrograph 红外成象器/摄谱仪IRPS, Infrared Photometer / Spectro- 红外光度计/分光计meterirregular cluster 不规则星团; 不规则星系团IRTF, NASA Infrared Telescope 〈IRTF〉美国宇航局红外Facility 望远镜IRTS, Infrared Telescope in Space 〈IRTS〉空间红外望远镜ISO, Infrared Space Observatory 〈ISO〉红外空间天文台isochrone method 等龄线法IUE, International Ultraviolet 〈IUE〉国际紫外探测器ExplorerJewel Box （NGC 4755 ）宝盒星团Jovian magnetosphere 木星磁层Jovian ring 木星环Jovian ringlet 木星细环Jovian seismology 木震学jovicentric orbit 木心轨道J-type star J 型星Juliet 天卫十一Jupiter-crossing asteroid 越木小行星Kalman filter 卡尔曼滤波器KAO, Kuiper Air-borne Observatory 〈柯伊伯〉机载望远镜Keck ⅠTelescope 凯克Ⅰ望远镜Keck ⅡTelescope 凯克Ⅱ望远镜Kuiper belt 柯伊伯带Kuiper-belt object 柯伊伯带天体Kuiper disk 柯伊伯盘LAMOST, Large Multi-Object Fibre 大型多天体分光望远镜Spectroscopic TelescopeLaplacian plane 拉普拉斯平面late cluster 晚型星系团LBT, Large Binocular Telescope 〈LBT〉大型双筒望远镜lead oxide vidicon 氧化铅光导摄象管Leo Triplet 狮子三重星系LEST, Large Earth-based Solar 〈LEST〉大型地基太阳望远镜Telescopelevel-Ⅰcivilization Ⅰ级文明level-Ⅱcivilization Ⅱ级文明level-Ⅲcivilization Ⅲ级文明Leverrier ring 勒威耶环Liapunov characteristic number 李雅普诺夫特征数（LCN ）light crown 轻冕玻璃light echo 回光light-gathering aperture 聚光孔径light pollution 光污染light sensation 光感line image sensor 线成象敏感器line locking 线锁line-ratio method 谱线比法Liner, low ionization nuclear 低电离核区emission-line regionline spread function 线扩散函数LMT, Large Millimeter Telescope 〈LMT〉大型毫米波望远镜local galaxy 局域星系local inertial frame 局域惯性架local inertial system 局域惯性系local object 局域天体local star 局域恒星look-up table （LUT ）对照表low-mass X-ray binary 小质量X 射线双星low-metallicity cluster 低金属度星团;低金属度星系团low-resolution spectrograph 低分辨摄谱仪low-resolution spectroscopy 低分辨分光low - z 小红移luminosity mass 光度质量luminosity segregation 光度层化luminous blue variable 高光度蓝变星lunar atmosphere 月球大气lunar chiaroscuro 月相图Lunar Prospector 〈月球勘探者〉Ly-α forest 莱曼-α 森林MACHO （massive compact halo 晕族大质量致密天体object ）Magellan 〈麦哲伦〉金星探测器Magellan Telescope 〈麦哲伦〉望远镜magnetic canopy 磁蓬magnetic cataclysmic variable 磁激变变星magnetic curve 磁变曲线magnetic obliquity 磁夹角magnetic period 磁变周期magnetic phase 磁变相位magnitude range 星等范围main asteroid belt 主小行星带main-belt asteroid 主带小行星main resonance 主共振main-sequence band 主序带Mars-crossing asteroid 越火小行星Mars Pathfinder 火星探路者mass loss rate 质量损失率mass segregation 质量层化Mayall Telescope 梅奥尔望远镜Mclntosh classification 麦金托什分类McMullan camera 麦克马伦电子照相机mean motion resonance 平均运动共振membership of cluster of galaxies 星系团成员membership of star cluster 星团成员merge 并合merger 并合星系; 并合恒星merging galaxy 并合星系merging star 并合恒星mesogranulation 中米粒组织mesogranule 中米粒metallicity 金属度metallicity gradient 金属度梯度metal-poor cluster 贫金属星团metal-rich cluster 富金属星团MGS, Mars Global Surveyor 火星环球勘测者micro-arcsec astrometry 微角秒天体测量microchannel electron multiplier 微通道电子倍增管microflare 微耀斑microgravitational lens 微引力透镜microgravitational lensing 微引力透镜效应microturbulent velocity 微湍速度millimeter-wave astronomy 毫米波天文millisecond pulsar 毫秒脉冲星minimum mass 质量下限minimum variance 最小方差mixed-polarity magnetic field 极性混合磁场MMT, Multiple-Mirror Telescope 多镜面望远镜moderate-resolution spectrograph 中分辨摄谱仪moderate-resolution spectroscopy 中分辨分光modified isochrone method 改进等龄线法molecular outflow 外向分子流molecular shock 分子激波monolithic-mirror telescope 单镜面望远镜moom 行星环卫星moon-crossing asteroid 越月小行星morphological astronomy 形态天文morphology segregation 形态层化MSSSO, Mount Stromlo and Siding 斯特朗洛山和赛丁泉天文台Spring Observatorymultichannel astrometric photometer 多通道天测光度计（MAP ）multi-object spectroscopy 多天体分光multiple-arc method 复弧法multiple redshift 多重红移multiple system 多重星系multi-wavelength astronomy 多波段天文multi-wavelength astrophysics 多波段天体物。

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arXiv:astro-ph/0103360v2 26 Apr 2001AccretionModelsofGamma-RayBurstsRameshNarayan1,TsviPiran2andPawanKumar31.Harvard-SmithsonianCenterforAstrophysics,Cambridge,MA02138,USA2.RacahInstituteofPhysics,HebrewUniversity,Jerusalem91904,Israel3.InstituteforAdvancedStudies,Princeton,NJ08540,USA.

ABSTRACTManymodelsofgamma-raybursts(GRBs)involveaccretionontoacompactobject,usuallyablackhole,atamassaccretionrateoforderafractionofasolarmasspersecond.IftheaccretiondiskislargerthanafewtensorhundredsofSchwarzschildradii,theaccretionwillproceedviaaconvection-dominatedaccretionﬂow(CDAF)inwhichmostofthematterescapestoinﬁnityratherthanfallingontotheblackhole.Modelsinvolvingthemergersofblackholewhitedwarfbinariesandblackholeheliumstarbinariesfallinthiscategory.ThesemodelsareunlikelytoproduceGRBssinceverylittlemassreachestheblackhole.Iftheaccretiondiskissmaller,thenaccretionwillproceedvianeutrinocoolinginaneutrino-dominatedaccretiondisk(NDAF)andmostofthemasswillreachthecenter.ModelsinvolvingthemergersofdoubleneutronstarbinariesandblackholeneutronstarbinariesfallinthiscategoryandarecapableofproducingbrightGRBs.IftheviscosityparameterαintheNDAFhasastandardvalue∼0.1,thesemergerscanexplainshortGRBswithdurationsunderasecond,buttheyareunlikelytoproducelongGRBswithdurationsoftensorhundredofseconds.Iftheaccretiondiskisfedbyfallbackofmaterialafterasupernovaexplosion,asinthecollapsarmodel,thenthetimescaleoftheburstisdeterminedbyfallback,notaccretion.SuchamodelcanproducelongGRBs.FallbackmodelsagainrequirethattheaccretionshouldproceedviaanNDAFratherthanaCDAFinorderforasigniﬁcantamountofmasstoreachtheblackhole.Thisconditionimposesanupperlimitontheradiusofinjectionofthegas.

1.IntroductionTheﬁreballmodel(seePiran1999,2000,forreviews)providesagoodunderstandingofconditionswithintheγ-ray-emittingandafterglow-emittingregionsofgamma-raybursts(GRBs).Accordingtothismodel,GRBsareproducedwhenrelativisticejectafroma“centralengine”aresloweddownbyinteractions,eitherwithanexternalmedium(theexternalshockmodel)oramongdiﬀerentlayerswithintheejectathemselves(theinternalshockmodel).Intheinteractionsthekineticenergyintheejectaisconvertedtorelativisticelectronswhichproducetheobserved–2–radiation.Amongthemanysuccessesofthemodelwenotetheobservationalconﬁrmationofrelativisticmotionintheafterglow(Frailetal.1997;Katz&Piran1997).

Despitethesuccessesoftheﬁreballmodel,thenatureofthecentralengineremainsamystery.Theproblemisthatthecentralengineishiddenfromview;noradiation(apartfromgravitationalradiationandneutrinosthatmaypossiblybedetectedinthedistantfuture)reachestheobserverdirectlyfromtheengine.Foranumberofso-called“longbursts”accuratepositionshavebeendeterminedthroughobservationsoftheirafterglows.Basedonthis,thereiscircumstantialevidencethattheseburstsareassociatedwithstar-formingregions(e.g.Bloometal.2000).ThereisnoinformationatpresentontheotherclassofGRBs,theso-called“shortbursts.”

Althoughwelackdirectevidenceonthenatureofthecentralengine,itisneverthelesswidelyacceptedthatGRBsaretheresultofcataclysmiceventsinvolvingeitherneutronstarsorstellar-massblackholes.Theargumentsinsupportofthishypothesisarestraightforward.(i)Sinceburstsradiatethebulkoftheirenergyintheγ-rayband,itseemslikelythatarelativisticobjectisbehindtheirproduction.(ii)Theenergybudget(∼1051erg)iscomparabletothekineticenergyofejectainasupernovaexplosion.(iii)Mostlongburstsarehighlyvariableingamma-rays,andsoaremanyshortbursts(Nakar&Piran,2001a).Inparticular,theratioofthetotaldurationofthebursttothevariabilitytimescaleislarge,fromwhichoneconcludesthatthegamma-raysmustbeproducedininternalshocks(Sari&Piran,1997).Akeyfeatureofinternalshocksisthattheobservedgamma-rayvariabilityreﬂectsthevariabilityintheactivityofthecentralengine(Sari&Piran,1997).Sincevariabilitytimescalesasshortasamillisecondareobserved,theenginemustcontainacompactobjectofnomorethanafewsolarmasses(otherwisethelight-crossingtimewouldexceedthevariabilitytime).

AninterestingcluetothenatureofGRBsisprovidedbythedurationsofbursts.Whilethefastestvariabilitytimescaleisunderamillisecond,burstdurationsareusuallyverymuchlonger.Longburstshavedurationsrangingfrom10–1000seconds,andevenshortburstshaveamediandurationofabout0.3seconds.Clearly,whateveristhephysicalmechanismbehindGRBproduction,itactsonamuchlongertimescalethanthefastestdynamicaltimeofthecentralengine.

Narayan,Paczy´nski&Piran(1992)suggestedthatthecentralengineinGRBsinvolvestheaccretionofmatterontoacompactstar,andthattheenergyintheburstisprovidedbythegravitationalenergyreleasedbytheaccretinggas.Insuchamodel,thedurationoftheburstissetbytheviscoustimescaleoftheaccretinggas.Inmostaccretionﬂows,theviscoustimeissigniﬁcantlylongerthanthedynamicaltime,andsotheaccretionmodelnaturallyexplainsthelargediﬀerencebetweenthedurationsofburstsandtheirfastestvariabilitytime.