Testing Quantum Gravity with LIGO and VIRGO
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Pulsar Physics Without Magnetars
Chin. J. Astron. Astrophys. Vol.8 (2008), Supplement, 213–218 Chinese Journal of (http://www.chjaa.org) Astronomy and Astrophysics Pulsar Physics without Magnetars Wolfgang Kundt Argelander-Institut f¨ur Astronomie der Universit¨at, Auf dem H¨ugel 71, D-53121 Bonn, Germany Abstract Magnetars, as defined some 15 yr ago, are inconsistent both with fundamental physics, and with the (by now) two observed upward jumps of the period derivatives (of two of them). Instead, the class of peculiar X-ray sources commonly interpreted as magnetars can be understood as the class of throttled pulsars, i.e. of pulsars strongly interacting with their CSM. This class is even expected to harbour the sources of the Cosmic Rays, and of all the (extraterrestrial) Gamma-Ray Bursts. Key words: magnetars — throttled pulsars — dying pulsars — cavity — low-mass disk 1 DEFINITION AND PROPERTIES OF THE MAGNETARS Magnetars have been defined by Duncan and Thompson some 15 years ago, as spinning, compact X-ray sources - probably neutron stars - which are powered by the slow decay of their strong magnetic field, of strength 1015 G near the surface, cf. (Duncan & Thompson 1992; Thompson & Duncan 1996). They are now thought to comprise the anomalous X-ray pulsars (AXPs), soft gamma-ray repeaters (SGRs), recurrent radio transients (RRATs), or ‘stammerers’, or ‘burpers’, and the ‘dim isolated neutron stars’ (DINSs), i.e. a large, fairly well defined subclass of all neutron stars, which has the following properties (Mereghetti et al. 2002): 1) They are isolated neutron stars, with spin periods P between 5 s and 12 s, and similar glitch behaviour to other neutron-star sources, which correlates with their X-ray bursting. -
Progenitors of Magnetars and Hyperaccreting Magnetized Disks
Feeding Compact Objects: Accretion on All Scales Proceedings IAU Symposium No. 290, 2012 c International Astronomical Union 2013 C. M. Zhang, T. Belloni, M. M´endez & S. N. Zhang, eds. doi:10.1017/S1743921312020340 Progenitors of Magnetars and Hyperaccreting Magnetized Disks Y. Xie1 andS.N.Zhang1,2 1 National Astronomical Observatories, Chinese Academy Of Sciences, Beijing, 100012, P. R.China email: [email protected] 2 Key Laboratory of Particle Astrophysics, Institute of High Energy Physics, Chinese Academy of Sciences, Beijing, 100012, P. R. China email: [email protected] Abstract. We propose that a magnetar could be formed during the core collapse of massive stars or coalescence of two normal neutron stars, through collecting and inheriting the magnetic fields magnified by hyperaccreting disk. After the magnetar is born, its dipole magnetic fields in turn have a major influence on the following accretion. The decay of its toroidal field can fuel the persistent X-ray luminosity of either an SGR or AXP; however the decay of only the poloidal field is insufficient to do so. Keywords. Magnetar, accretion disk, gamma ray bursts, magnetic fields 1. Introduction Neutron stars (NSs) with magnetic field B stronger than the quantum critical value, 2 3 13 Bcr = m c /e ≈ 4.4 × 10 G, are called magnetars. Observationally, a magnetar may appear as a soft gamma-ray repeater (SGR) or an anomalous X-ray pulsar (AXP). It is believed that its persistent X-ray luminosity is powered by the consumption of their B- decay energy (e.g. Duncan & Thompson 1992; Paczynski 1992). However, the formation of strong B of a magnetar remains unresolved, and the possible explanations are divided into three classes: (i) B is generated by a convective dynamo (Duncan & Thompson 1992); (ii) B is essentially of fossil origin (e.g. -
Mass – Luminosity Relation for Massive Stars
MASS – LUMINOSITY RELATION FOR MASSIVE STARS Within the Eddington model β ≡ Pg/P = const, and a star is an n = 3 polytrope. Large mass stars have small β, and hence are dominated by radiation pressure, and the opacity in them is dominated by electron scattering. Let us consider the outer part of such a star assuming it is in a radiative equilibrium. We have the equation of hydrostatic equilibrium: dP GM = − r ρ, (s2.1) dr r2 and the equation of radiative equilibrium dP κρL r = − r . (s2.2) dr 4πcr2 Dividing these equations side by side we obtain dP κL r = r , (s2.3) dP 4πcGMr Near the stellar surface we have Mr ≈ M and Lr ≈ L, and adopting κ ≈ κe = const, we may integrate equation (s2.3) to obtain κeLr Pr − Pr,0 = (P − P0) , (s2.4) 4πcGMr 4 where Pr,0 = P0 = aTeff /6 is the radiation pressure at the surface, i.e. at τ = 0, where the gas pressure Pg = 0. At a modest depth below stellar surface pressure is much larger than P0, we may neglect the integration constants in (s2.4) to obtain P κ L L 1+ X M⊙ L 1 − β = r = e = = , (s2.5) P 4πcGM LEdd 65300 L⊙ M where LEdd is the Eddington luminosity. Equation (s2.5) gives a relation between stellar mass, luminosity and β. The Eddington model gives a relation between stellar mass and β : 1 2 M 18.1 (1 − β) / = 2 2 , (s2.6) M⊙ µ β where µ is a mean molecular weight in units of mass of a hydrogen atom, µ−1 =2X +0.75Y +0.5Z, and X,Y,Z, are the hydrogen, helium, and heavy element abundance by mass fraction. -
Astronomy 112: the Physics of Stars Class 10 Notes: Applications and Extensions of Polytropes in the Last Class We Saw That Poly
Astronomy 112: The Physics of Stars Class 10 Notes: Applications and Extensions of Polytropes In the last class we saw that polytropes are a simple approximation to a full solution to the stellar structure equations, which, despite their simplicity, yield important physical insight. This is particularly true for certain types of stars, such as white dwarfs and very low mass stars. In today’s class we will explore further extensions and applications of simple polytropic models, which we can in turn use to generate our first realistic models of typical main sequence stars. I. The Binding Energy of Polytropes We will begin by showing that any realistic model must have n < 5. Of course we already determined that solutions to the Lane-Emden equation reach Θ = 0 at finite ξ only for n < 5, which already hints that n ≥ 5 is a problem. Nonetheless, we have not shown that, as a matter of physical principle, models of this sort are unacceptable for stars. To demonstrate this, we will prove a generally useful result about the energy content of polytropic stars. As a preliminary to this, we will write down the polytropic relation (n+1)/n P = KP ρ in a slightly different form. Consider a polytropic star, and imagine moving down within it to the point where the pressure is larger by a small amount dP . The corre- sponding change in density dρ obeys n + 1 dP = K ρ1/ndρ. P n Similarly, since P = K ρ1/n, ρ P it follows that P ! K 1 dP ! d = P ρ(1−n)/ndρ = ρ n n + 1 ρ We can regard this relation as telling us how much the ratio of pressure to density changes when we move through a star by an amount such that the pressure alone changes by dP . -
Ultraluminous X-Ray Pulsars
Ultraluminous X-ray pulsars: the high-B interpretation Juri Poutanen (University of Turku, Finland; IKI, Moscow) Sergey Tsygankov, Alexander Mushtukov, Valery Suleimanov, Alexander Lutovinov, Anna Chashkina, Pavel Abolmasov Plan: • ULX before 2014 • A short history of ULXPs • Supercritical accretion onto a magnetized neutron star • Large or low B? Propeller? • Beamed? • Winds? Models for ULX (before 2014) • Super-Eddington accretion onto a stellar-mass black hole (e.g. King 2001, Begelman et al. 2006, Poutanen et al. 2007) • Sub-Eddington accretion onto intermediate mass black holes (Colbert & Mushotzky 2001) • Young rotation-powered pulsar (Medvedev & Poutanen 2013) Super-Eddington accretion • Slim disk models Accretion rate is large, but most of the released energy is advected towards the BH. ˙ ˙ M (r) =M 0 • Super-disks with winds Accretion rate is large, but most of the mass is blown away by radiation. Only the Eddington rate goes to the BH. ˙ ˙ r M (r) =M 0 Rsph ˙ ˙ M (rin ) =M Edd M˙ R R 0 - spherization radius sph ≈ in ˙ M Edd Shakura & Sunyaev 1973 Super-Eddington accretion • Slim disk models Accretion rate is large, but most of the released energy is advected towards the BH. ˙ ˙ M (r) =M 0 • Super-disks with winds Accretion rate is large, but most of the mass is blown away by radiation. Only the Eddington rate goes to the BH. ! ! r M (r) =M 0 Rsph ! ! M(rin ) =M Edd M! R ≈ R 0 -spherization radius sph S ! M Edd Ohsuga et al. 2005 Super-Eddington accretion Super-disks with winds and advection Accretion rate is large, a lot of mass is blown away by radiation (using fraction εW of available radiative flux), but still a significant fraction goes to the BH. -
MHD Accretion Disk Winds: the Key to AGN Phenomenology?
galaxies Review MHD Accretion Disk Winds: The Key to AGN Phenomenology? Demosthenes Kazanas NASA/Goddard Space Flight Center, 8800 Greenbelt Rd, Greenbelt, MD 20771, USA; [email protected]; Tel.: +1-301-286-7680 Received: 8 November 2018; Accepted: 4 January 2019; Published: 10 January 2019 Abstract: Accretion disks are the structures which mediate the conversion of the kinetic energy of plasma accreting onto a compact object (assumed here to be a black hole) into the observed radiation, in the process of removing the plasma’s angular momentum so that it can accrete onto the black hole. There has been mounting evidence that these structures are accompanied by winds whose extent spans a large number of decades in radius. Most importantly, it was found that in order to satisfy the winds’ observational constraints, their mass flux must increase with the distance from the accreting object; therefore, the mass accretion rate on the disk must decrease with the distance from the gravitating object, with most mass available for accretion expelled before reaching the gravitating object’s vicinity. This reduction in mass flux with radius leads to accretion disk properties that can account naturally for the AGN relative luminosities of their Optical-UV and X-ray components in terms of a single parameter, the dimensionless mass accretion rate. Because this critical parameter is the dimensionless mass accretion rate, it is argued that these models are applicable to accreting black holes across the mass scale, from galactic to extragalactic. Keywords: accretion disks; MHD winds; accreting black holes 1. Introduction-Accretion Disk Phenomenology Accretion disks are the generic structures associated with compact objects (considered to be black holes in this note) powered by matter accretion onto them. -
Chapter 9 Accretion Disks for Beginners
Chapter 9 Accretion Disks for Beginners When the gas being accreted has high angular momentum, it generally forms an accretion disk. If the gas conserves angular momentum but is free to radiate energy, it will lose energy until it is on a circular orbit of radius 2 Rc = j /(GM), where j is the specific angular momentum of the gas, and M is the mass of the accreting compact object. The gas will only be able to move inward from this radius if it disposes of part of its angular momentum. In an accretion disk, angular momentum is transferred by viscous torques from the inner regions of the disk to the outer regions. The importance of accretion disks was first realized in the study of binary stellar systems. Suppose that a compact object of mass Mc and a ‘normal’ star of mass Ms are separated by a distance a. The normal star (a main- sequence star, a giant, or a supergiant) is the source of the accreted matter, and the compact object is the body on which the matter accretes. If the two bodies are on circular orbits about their center of mass, their angular velocity will be 1/2 ~ G(Mc + Ms) Ω= 3 eˆ , (9.1) " a # wheree ˆ is the unit vector normal to the orbital plane. In a frame of reference that is corotating with the two stars, the equation of momentum conservation has the form ∂~u 1 +(~u · ∇~ )~u = − ∇~ P − 2Ω~ × ~u − ∇~ Φ , (9.2) ∂t ρ R 89 90 CHAPTER 9. ACCRETION DISKS FOR BEGINNERS Figure 9.1: Sections in the orbital plane of equipotential surfaces, for a binary with q = 0.2. -
Exceeding the Eddington Limit
The Fate of the Most Massive Stars ASP Conference Series, Vol. 332, 2005 Roberta M. Humphreys and Krzysztof Z. Stanek Exceeding the Eddington Limit Nir J. Shaviv Racah Institute of Physics, Hebrew University of Jerusalem, Jerusalem 91904, Israel Abstract. We review the current theory describing the existence of steady state super-Eddington atmospheres. The key to the understanding of these at- mospheres is the existence of a porous layer responsible for a reduced e®ective opacity. We show how porosity arises from radiative-hydrodynamic instabili- ties and why the ensuing inhomogeneities reduce the e®ective opacity. We then discuss the appearance of these atmospheres. In particular, one of their funda- mental characteristic is the continuum driven acceleration of a thick wind from regions where the inhomogeneities become transparent. We end by discussing the role that these atmospheres play in the evolution of massive stars. 1. Introduction The notion that astrophysical objects can remain in a super-Eddington steady state for considerable amounts of time, contradicts the common wisdom usually invoked in which the classical Eddington limit, Edd, cannot be exceeded in a steady state because no hydrostatic solution canLthen exist. In other words, if objects do pass Edd, they are expected to be highly dynamic. A huge mass loss should ensueLsince large parts of their atmospheres are then gravitationally unbound and should therefore be expelled. ´ Carinae, and classical novae for that matter, stand in contrast to the above notion. ´ Car was super-Eddington (hereafter SED) during its 20 year long eruption (Davidson and Humphreys 1997). Yet, its observed mass loss and velocity are inconsistent with any standard solution (Shaviv 2000). -
Timescales of the Solar Protoplanetary Disk 233
Russell et al.: Timescales of the Solar Protoplanetary Disk 233 Timescales of the Solar Protoplanetary Disk Sara S. Russell The Natural History Museum Lee Hartmann Harvard-Smithsonian Center for Astrophysics Jeff Cuzzi NASA Ames Research Center Alexander N. Krot Hawai‘i Institute of Geophysics and Planetology Matthieu Gounelle CSNSM–Université Paris XI Stu Weidenschilling Planetary Science Institute We summarize geochemical, astronomical, and theoretical constraints on the lifetime of the protoplanetary disk. Absolute Pb-Pb-isotopic dating of CAIs in CV chondrites (4567.2 ± 0.7 m.y.) and chondrules in CV (4566.7 ± 1 m.y.), CR (4564.7 ± 0.6 m.y.), and CB (4562.7 ± 0.5 m.y.) chondrites, and relative Al-Mg chronology of CAIs and chondrules in primitive chon- drites, suggest that high-temperature nebular processes, such as CAI and chondrule formation, lasted for about 3–5 m.y. Astronomical observations of the disks of low-mass, pre-main-se- quence stars suggest that disk lifetimes are about 3–5 m.y.; there are only few young stellar objects that survive with strong dust emission and gas accretion to ages of 10 m.y. These con- straints are generally consistent with dynamical modeling of solid particles in the protoplanetary disk, if rapid accretion of solids into bodies large enough to resist orbital decay and turbulent diffusion are taken into account. 1. INTRODUCTION chondritic meteorites — chondrules, refractory inclusions [Ca,Al-rich inclusions (CAIs) and amoeboid olivine ag- Both geochemical and astronomical techniques can be gregates (AOAs)], Fe,Ni-metal grains, and fine-grained ma- applied to constrain the age of the solar system and the trices — are largely crystalline, and were formed from chronology of early solar system processes (e.g., Podosek thermal processing of these grains and from condensation and Cassen, 1994). -
Ast 777: Star and Planet Formation Accretion Disks
Ast 777: Star and Planet Formation Jonathan Williams, University of Hawaii Accretion disks https://www.universetoday.com/135153/first-detailed-image-accretion-disk-around-young-star/ The main phases of star formation https://www.americanscientist.org/sites/americanscientist.org/files/2005223144527_306.pdf Star grows by accretion through a disk https://www.americanscientist.org/sites/americanscientist.org/files/2005223144527_306.pdf Accretion continues (at a low rate) through the CTTS phase https://www.americanscientist.org/sites/americanscientist.org/files/2005223144527_306.pdf The mass of a star is, by far, the dominant parameter that determines its fate. The primary question for star formation is how stars reach their ultimate mass. Mass infall rate for a BE core that collapses A strikingly simple equation! in a free-fall time isothermal sound speed calculated for T=10K This is very high — if sustained, a solar mass star would form in 0.25 Myr… …but it can’t all fall into the center Conservation of angular momentum Bonnor-Ebert sphere, R ~ 0.2pc ~ 6x1017cm 11 Protostar, R ~ 3R⊙ ~ 2x10 cm => gravitational collapse results in a compression of spatial scales > 6 orders of magnitude! => spin up! https://figureskatingmargotzenaadeline.weebly.com/angular-momentum.html Conservation of angular momentum Angular momentum Moment of inertia J is conserved Galactic rotation provides a lower limit to core rotation (R⊙=8.2kpc, V⊙=220 km/s) Upper limit to rotation period of the star So how do stars grow? • Binaries (~50% of field stars are in multiples) -
The Unique Potential of Extreme Mass-Ratio Inspirals for Gravitational-Wave Astronomy 1, 2 3, 4 Christopher P
Main Thematic Area: Formation and Evolution of Compact Objects Secondary Thematic Areas: Cosmology and Fundamental Physics, Galaxy Evolution Astro2020 Science White Paper: The unique potential of extreme mass-ratio inspirals for gravitational-wave astronomy 1, 2 3, 4 Christopher P. L. Berry, ∗ Scott A. Hughes, Carlos F. Sopuerta, Alvin J. K. Chua,5 Anna Heffernan,6 Kelly Holley-Bockelmann,7 Deyan P. Mihaylov,8 M. Coleman Miller,9 and Alberto Sesana10, 11 1CIERA, Northwestern University, 2145 Sheridan Road, Evanston, IL 60208, USA 2Department of Physics and MIT Kavli Institute, Massachusetts Institute of Technology, Cambridge, MA 02139, USA 3Institut de Ciencies` de l’Espai (ICE, CSIC), Campus UAB, Carrer de Can Magrans s/n, 08193 Cerdanyola del Valles,` Spain 4Institut d’Estudis Espacials de Catalunya (IEEC), Edifici Nexus I, Carrer del Gran Capita` 2-4, despatx 201, 08034 Barcelona, Spain 5Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, CA 91109, USA 6School of Mathematics and Statistics, University College Dublin, Belfield, Dublin 4, Ireland 7Department of Physics and Astronomy, Vanderbilt University and Fisk University, Nashville, TN 37235, USA 8Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge, CB3 0HA, UK 9Department of Astronomy and Joint Space-Science Institute, University of Maryland, College Park, MD 20742-2421, USA 10School of Physics & Astronomy and Institute for Gravitational Wave Astronomy, University of Birmingham, Edgbaston, Birmingham B15 2TT, UK 11Universita` di Milano Bicocca, Dipartimento di Fisica G. Occhialini, Piazza della Scienza 3, I-20126, Milano, Italy (Dated: March 7, 2019) The inspiral of a stellar-mass compact object into a massive ( 104–107M ) black hole ∼ produces an intricate gravitational-wave signal. -
Arxiv:1707.03021V4 [Gr-Qc] 8 Oct 2017
The observational evidence for horizons: from echoes to precision gravitational-wave physics Vitor Cardoso1,2 and Paolo Pani3,1 1CENTRA, Departamento de F´ısica,Instituto Superior T´ecnico, Universidade de Lisboa, Avenida Rovisco Pais 1, 1049 Lisboa, Portugal 2Perimeter Institute for Theoretical Physics, 31 Caroline Street North Waterloo, Ontario N2L 2Y5, Canada 3Dipartimento di Fisica, \Sapienza" Universit`adi Roma & Sezione INFN Roma1, Piazzale Aldo Moro 5, 00185, Roma, Italy October 10, 2017 Abstract The existence of black holes and of spacetime singularities is a fun- damental issue in science. Despite this, observations supporting their existence are scarce, and their interpretation unclear. We overview how strong a case for black holes has been made in the last few decades, and how well observations adjust to this paradigm. Unsurprisingly, we con- clude that observational proof for black holes is impossible to come by. However, just like Popper's black swan, alternatives can be ruled out or confirmed to exist with a single observation. These observations are within reach. In the next few years and decades, we will enter the era of precision gravitational-wave physics with more sensitive detectors. Just as accelera- tors require larger and larger energies to probe smaller and smaller scales, arXiv:1707.03021v4 [gr-qc] 8 Oct 2017 more sensitive gravitational-wave detectors will be probing regions closer and closer to the horizon, potentially reaching Planck scales and beyond. What may be there, lurking? 1 Contents 1 Introduction 3 2 Setting the stage: escape cone, photospheres, quasinormal modes, and tidal effects 6 2.1 Geodesics . .6 2.2 Escape trajectories .