DOCSLIB.ORG

7 Stellar Winds & Supernovae Remnants

Total Page：16

File Type：pdf, Size：1020Kb

Download full-text PDF Read full-text
7 Stellar Winds & Supernovae Remnants 36 Physics 426 Notes Spring 2009 7 Stellar Winds & Supernovae Remnants Whether or not a particular flow satisfies this condition depends on the starting conditions, such as with what Now let’s look at some supersonic flow situations velocity and temperature it left the stellar surface, and which involve shocks. You remember that we worked also what the boundary conditions at large distances with smooth transonic flows last term – for instance are. If it does not start in such a way to satisfy this the solar wind. These are rare; it’s very easy to form condition, it either stays subsonic (corresponding to fi- shocks when supersonic flows are decelerated or bent nite pressure at infinity), or cannot establish a steady by their surroundings. In this chapter we’ll work with flow. 2-1/2 examples: stellar wind shocks, and two types of Further, the solution beyond the sonic point depends supernova remnants. • on the temperature structure of the wind. The only so- 7.1 Stellar winds and the surrounding ISM lutions with dv/dr > 0 for r > rs are those for which 2 cs(r) drops off more slowly than 1/r; it is only these We worked with smooth, transonic stellar wind flow for which the right-hand side stays positive. In the case last term ... and found a solution in which the outer 2 of an isothermal wind, with cs = constant, (7.3) can be regions of the wind flow are supersonic. But this can’t solved in the limit r rs: carry on forever. At some point the wind flow must run " into the ambient ISM. We expect the deceration to lead v2(r) 4c2 ln r + constant (7.5) # s to an outer shock in the wind; what are the details? By way of review, I’ll repeat the basics of the inner-wind Thus, the wind will be supersonic, by a factor of a few, as r . The question of how the solar wind man- solution here. → ∞ ages to stay nearly isothermal is not solved; it is prob- 7.1.1 The basic solution ably due to energy transport by some sort of waves (MHD or plasma waves, for instance) which are gen- Mass conservation in a steady, spherical flow is • erated in the photosphere and damped somewhere far ρvr2 = constant; or, out in the wind. 1 dρ 1 dv 2 + + = 0 (7.1) ρ dr v dr r 7.1.2 The outer shock while the momentum equation becomes in this case The pressure in this supersonic wind is dropping with (noting that gravity from the central star is important), radius (since ρ 1/vr2, with v being only slowly ∝ varying; thus p ρT 1/r2, approximately, in dv dp GM ∝ ∝ ρv + = ρ (7.2) an isothermal wind). Therefore, when the wind pres- dr dr − r2 sure is close to the ISM pressure, the wind must slow 2 Writing dp/dr = csdρ/dr, these two equations com- down. At this outer boundary, we expect some sort of bine to give the basic wind equation, shock transition, since the wind is supersonic. Past this c2 dv 2c2 GM shock, the hot, shocked wind-gas will expand into the v s = s (7.3) ISM (at about its own sound speed, to start); as long − v dr r − r2 ! " as this expansion is supersonic relative to the ISM, the This does not have analytic solutions over the whole expanding hot gas will drive a “snowplowed” shell of range of r. However, we can learn quite a bit about the ISM, and a second shock, out into the ISM. nature of the solutions simply by inspection of (7.3), as A cartoon of this region, at some point in time, would follows. be that in Figure 7.1. Let region “a” be the wind; S1 First, the left hand side contains a zero, at v2 = c2. be the inner shock; region “b” be the wind-gas which • s If we want to consider well-behaved flows, that is to has been through the shock; C be the contact sur- say those in which the derivative dv/dr does not blow face between the wind and the ISM; region “c” be the up, then the right hand side of (7.3) must go to zero at shocked ISM; and S2 be the outer shock (moving into the same point. This defines the condition that must be the ISM). We expect S1 to be an adiabatic shock (since met at the sonic point: the wind is probably hot and low density, and thus will have a long cooling time); region “b” will contain hot, 2 2 GM 2 v = c at r = rs = (7.4) 3 mvwind 7 s 2c2 shocked wind, with Tb several 10 K s ∼ 16 kB ∼ × Physics 426 Notes Spring 2009 37 1 (noting that m = 2 mp is the mean mass per particle if bit of algebra tells us that α = 3/5 is the allowed solu- region “b” is fully ionized). The outer shock will prob- tion, and we can also find an expression for A. Thus, ably be isothermal, since the ISM is denser and cooler the solution is than the wind. Thus, the shocked ISM will be in a 1/5 E˙ thin shell, containing all of the original ISM that lay R(t) = 0.76 wind t3/5 (7.9) between S2 and the star. % ρo & S 2 and C ˙ 1/5 dR Ewind −2/5 S 1 v(t) = = 0.46 t (7.10) dt % ρo & a Thus, the shell decelerates with time, as it ought to if R(t) b it is picking up more and more ISM. One can then use c d these, and (7.6), to work out the pressure acting on the outer shell: Figure 7.1 Cartoon of the structure of a stellar wind, and 7 2 −4/5 its interaction with the ISM. Left: the shock structure within pb(t) = A ρot (7.11) the wind. Right: The outer shell of dense, snowplowed ISM. 25 From Dyson & Williams figures 7.3 and 7.4. so that the outer pressure drops with time. Finally, one can also work out the kinetic energy of the shell, which To start, consider the position of this shell, R(t), as a must give the energy input to the general ISM from the function of time. We start with a force equation. The wind. This turns out to be mass is the ISM shell is the mass that was originally 2 2π dR within R(t): 4π ρ R3(t). The force acting on the shell R3ρ 0.2E˙ t (7.12) 3 o 3 o dt # wind is just the pressure behind it, which drives it outward: ! " so that about 20% of the wind energy goes to the ISM. d 4π 3 dR 2 (The rest goes to heating the bubble, and to “pdV” R ρo = 4πR pb (7.6) dt 3 dt work). ! " if pb is the pressure, in the outer part of region “b”, 7.1.3 What about the inner shock? which acts on the shell. Next, we need an equation The location of the inner shock, S1, is determined by a for pb(t). The energy within region “b” – which we will take to be nearly all of the volume within R(t) combination of the jump conditions, applied at S1, and – is 2πp R3(t); the rate of change of this energy the pressure at the outside of the region, pb, as follows. # b is given by the difference between the input (from the The jump conditions, at S1, are ρsb = 4ρsa, if “sb” • wind) and the “pdV” work done by the expansion: and “sa” subscripts refer to postshock (region “b”) and preshock (region “a”), respectively. Also, vsb = vsa/4 d 3 ˙ d 4π 3 – here we assume 1, the strong shock limit, 2πR pb = Ewind pb R (7.7) M " dt − dt 3 and also an adiabatic shock. At the shock, momentum ! " # $ conservation1 tells us that ρv2 + p is conserved; thus, if E˙ wind is the energy input rate from the wind, as- 3 2 3 2 sumed constant. These two equations can be com- p = ρsav + psa ρsav sb 4 sa # 4 sa bined, for instance by eliminating pb, to get where we have used the fact that v c in the wind sa " s 3 2 3 ˙ region. 4 d R 3 dR d R 2 dR 3 Ewind aR 3 + bR 2 + cR = dt dt dt dt 2π ρo In region “b”, where we can ignore gravity (com- ! " • 3 (7.8) pared to the internal energy, 2 p), the momentum equa- where a, b, c are numerical constants (that you’ll find in tion is dv dp the homework). Noting that each term on the left hand ρv + 0 side has the same dependence on R and t ( R5t−3), dr dr # ∼ we can try a power law solution, R(t) = Atα. A small 1Check back to chapter 4 from last term, P425; equation 4.4. 38 Physics 426 Notes Spring 2009 and, since v c in most of region “b”, ρ con- very high temperature and pressure, driving an expan- ) s # stant, and we have 1 ρv2 + p constant. Now – sion. This expansion will be very supersonic, setting 2 # we know that v drops from vsb = vsa/4, at S1, to up a spherical shock wave moving into the surround- v t−2/5 c at S (from the shock conditions). ings and sweeping up gas as it goes.
Load more

Recommended publications
  • Chapter 3 – Astrospheres, Stellar Winds, and the Interstellar Medium
    3 Astrospheres, stellar winds, and the interstellar medium BRIAN E. WOOD AND JEFFREY L. LINSKY 3.1 The spatial extent of the solar wind While orbiting the Sun, the Earth and the other planers in our solar system are ex posed to an outflow of plasma from the Sun. This is the solar wi nd, which largely determines the particle environment of the planets, as well as their magnetospheric properties. Long-tenn exposure Lo the solar wind may have had signjficant cffectc; on planetary atmospheres in some cases. Mars is a particularly interesting example, as the Martian atmosphere was once much thicker than it is today, and the erosive effects of solar wind exposure are commonly believed to be responsible for itc; depletion (see Ch. 8; also Jakosky et al.. 1994; Atreya et al., 2013). The solar wind does not expand indefinitely. Eventually it runs into the inter­ stellar medium (ISM), the ex tre mely low-particle-density e nvironment that exists in between the stars. Our Sun is moving relative to the ISM that surrounds the planetary system, so we see a fl ow of interstellar matter in the heliocentric rest frame, coming from the direction of the constellation Ophiuchus. The interaction between the solar wind and the ISM flow determines the large-scale structure of our he liosphere. which basically defines the extent of the solar wind's reach into our Galactic environment. Other stars are naturally surrounded by their own "astro­ spheres" (allernatively "asterospheres"') defined by the strength of their own stellar winds, the nature of the ISM in their Galactic neighborhoods, and their rnlative motion.
    [Show full text]
  • Lecture 17: the Solar Wind O Topics to Be Covered
    Lecture 17: The solar wind o Topics to be covered: o Solar wind o Inteplanetary magnetic field The solar wind o Biermann (1951) noticed that many comets showed excess ionization and abrupt changes in the outflow of material in their tails - is this due to a solar wind? o Assumed comet orbit perpendicular to line-of-sight (vperp) and tail at angle ! => tan! = vperp/vr o From observations, tan ! ~ 0.074 o But vperp is a projection of vorbit -1 => vperp = vorbit sin ! ~ 33 km s -1 o From 600 comets, vr ~ 450 km s . o See Uni. New Hampshire course (Physics 954) for further details: http://www-ssg.sr.unh.edu/Physics954/Syllabus.html The solar wind o STEREO satellite image sequences of comet tail buffeting and disconnection. Parker’s solar wind o Parker (1958) assumed that the outflow from the Sun is steady, spherically symmetric and isothermal. o As PSun>>PISM => must drive a flow. o Chapman (1957) considered corona to be in hydrostatic equibrium: dP = −ρg dr dP GMS ρ + 2 = 0 Eqn. 1 dr r o If first term >> than second €=> produces an outflow: € dP GM ρ dv + S + ρ = 0 Eqn. 2 dr r2 dt o This is the equation for a steadily expanding solar/stellar wind. € Parker’s solar wind (cont.) dv dv dr dv dP GM ρ dv o As, = = v => + S + ρv = 0 dt dr dt dr dr r2 dr dv 1 dP GM or v + + S = 0 Eqn. 3 dr ρ dr r2 € € o Called the momentum equation. € o Eqn. 3 describes acceleration (1st term) of the gas due to a pressure gradient (2nd term) and gravity (3rd term).
    [Show full text]
  • 1 Stellar Winds and Magnetic Fields
    1 Stellar winds and magnetic fields by Viggo Hansteen The solar wind is responsible for maintaining the heliosphere, and for being the driving agent in the magnetospheres of the planets but also for being the agent by which the Sun has shed angular momentum during the eons since its formation. We can assume that other cool stars with active corona have similar winds that play the same roles for those stars (1). The decade since the launch of SOHO has seen a large change in our understanding of the themrally driven winds such as the solar wind, due to both theoretical, computational, and observational advances. Recent solar wind models are characterized by low coronal electron temperatures while proton, α-particle, and minor ion temperatures are expected to be quite high and perhaps anisotropic. This entails that the electric field is relatively unimportant and that one quite naturally is able to obtain solar wind outflow that has both a high asymptotic flow speed and maintaining a low mass flux. In this chapter we will explain why these changes have come about and outline the questions now facing thermal wind astrophysicists. The progress we have seen in the last decade is largely due observations made with instruments onboard Ulysses (2) and SOHO (3). These observations have spawned a new understanding of solar wind energetics, and the consideration of the chromosphere, corona, and solar wind as a unified system. We will begin by giving our own, highly biased, \pocket history" of solar wind theory highlighting the problems that had to be resolved in order to make the original Parker formulation of thermally driven winds conform with observational results.
    [Show full text]
  • GEORGE HERBIG and Early Stellar Evolution
    GEORGE HERBIG and Early Stellar Evolution Bo Reipurth Institute for Astronomy Special Publications No. 1 George Herbig in 1960 —————————————————————– GEORGE HERBIG and Early Stellar Evolution —————————————————————– Bo Reipurth Institute for Astronomy University of Hawaii at Manoa 640 North Aohoku Place Hilo, HI 96720 USA . Dedicated to Hannelore Herbig c 2016 by Bo Reipurth Version 1.0 – April 19, 2016 Cover Image: The HH 24 complex in the Lynds 1630 cloud in Orion was discov- ered by Herbig and Kuhi in 1963. This near-infrared HST image shows several collimated Herbig-Haro jets emanating from an embedded multiple system of T Tauri stars. Courtesy Space Telescope Science Institute. This book can be referenced as follows: Reipurth, B. 2016, http://ifa.hawaii.edu/SP1 i FOREWORD I first learned about George Herbig’s work when I was a teenager. I grew up in Denmark in the 1950s, a time when Europe was healing the wounds after the ravages of the Second World War. Already at the age of 7 I had fallen in love with astronomy, but information was very hard to come by in those days, so I scraped together what I could, mainly relying on the local library. At some point I was introduced to the magazine Sky and Telescope, and soon invested my pocket money in a subscription. Every month I would sit at our dining room table with a dictionary and work my way through the latest issue. In one issue I read about Herbig-Haro objects, and I was completely mesmerized that these objects could be signposts of the formation of stars, and I dreamt about some day being able to contribute to this field of study.
    [Show full text]
  • Stellar Wind and Neutron Stars in Massive Binaries
    FISICPAC 2018 IOP Publishing IOP Conf. Series: Journal of Physics: Conf. Series 1258 (2019) 012021 doi:10.1088/1742-6596/1258/1/012021 Accretion flows around compact objects: stellar wind and neutron stars in massive binaries Antonios Manousakis Department of Applied Physics & Astronomy, University of Sharjah, Sharjah, UAE Sharjah Academy of Astronomy, Space Sciences, and Technology (SAASST), Sharjah, UAE E-mail: [email protected] Abstract. Strong winds from massive stars are a topic of interest to a wide range of astrophysical fields, e.g., stellar evolution, X-ray binaries or the evolution of galaxies. Massive stars, as heavy as 10 to 20 solar masses, generate dense fast outflows that continuously supply their surroundings with metals, kinetic energy and ionizing radiation, triggering star formation and driving the chemical enrichment and evolution of Galaxies. The amount of mass lost through the emission of these winds dramatically affects the evolution of the star. Depending on the mass left at the end of its life, a massive star can explode as a supernova and/or produce a gamma-ray burst, which are among the most powerful probes of cosmic evolution. The study of massive star winds is thus a truly multidisciplinary field and has a wide impact on different areas of astronomy. In High-Mass X-ray Binaries the presence of an accreting compact object on the one side allows to infer wind parameters from studies of the varying properties of the emitted X-rays; but on the other side the accretors gravity and ionizing radiation can strongly influence the wind flow. As the compact objects, either a neutron star or a black hole, moves through the dense stellar wind, it accretes material from its donor star and therefore emit X-rays.
    [Show full text]
  • Mass-Loss Predictions for O and B Stars As a Function of Metallicity
    A&A 369, 574–588 (2001) Astronomy DOI: 10.1051/0004-6361:20010127 & c ESO 2001 Astrophysics Mass-loss predictions for O and B stars as a function of metallicity Jorick S. Vink1,A.deKoter2, and H. J. G. L. M. Lamers1,3 1 Astronomical Institute, Utrecht University, PO Box 80000, 3508 TA Utrecht, The Netherlands 2 Astronomical Institute “Anton Pannekoek”, University of Amsterdam, Kruislaan 403, 1098 SJ Amsterdam, The Netherlands 3 SRON Laboratory for Space Research, Sorbonnelaan 2, 3584 CA Utrecht, The Netherlands Received 24 July 2000 / Accepted 17 January 2001 Abstract. We have calculated a grid of massive star wind models and mass-loss rates for a wide range of metal abundances between 1/100 ≤ Z/Z ≤ 10. The calculation of this grid completes the Vink et al. (2000) mass-loss recipe with an additional parameter Z. We have found that the exponent of the power law dependence of mass loss 0.85 p vs. metallicity is constant in the range between 1/30 ≤ Z/Z ≤ 3. The mass-loss rate scales as M˙ ∝ Z v∞ with p = −1.23 for stars with Teff ∼> 25 000 K, and p = −1.60 for the B supergiants with Teff ∼< 25 000 K. 0.13 Taking also into account the metallicity dependence of v∞, using the power law dependence v∞ ∝ Z from Leitherer et al. (1992), the overall result of mass loss as a function of metallicity can be represented by M˙ ∝ Z0.69 0.64 for stars with Teff ∼> 25 000 K, and M˙ ∝ Z for B supergiants with Teff ∼< 25 000 K.
    [Show full text]
  • Simulating Stellar Winds in AMUSE Edwin Van Der Helm1, Martha I
    A&A 625, A85 (2019) Astronomy https://doi.org/10.1051/0004-6361/201732020 & c ESO 2019 Astrophysics Simulating stellar winds in AMUSE Edwin van der Helm1, Martha I. Saladino2,1, Simon Portegies Zwart1, and Onno Pols2 1 Leiden Observatory, Leiden University, PO Box 9513, 2300 RA Leiden, The Netherlands e-mail: [email protected] 2 Department of Astrophysics/IMAPP, Radboud University, PO Box 9010, 6500 GL Nijmegen, The Netherlands Received 30 September 2017 / Accepted 24 March 2019 ABSTRACT Aims. We present stellar_wind.py, a module that provides multiple methods of simulating stellar winds using smoothed particle hydrodynamics codes (SPH) within the astrophysical multipurpose software environment (amuse) framework. Methods. The module currently includes three ways of simulating stellar winds: With the simple wind mode, we create SPH wind particles in a spherically symmetric shell for which the inner boundary is located at the radius of the star. We inject the wind particles with a velocity equal to their terminal velocity. The accelerating wind mode is similar, but with this method particles can be injected with a lower initial velocity than the terminal velocity and they are accelerated away from the star according to an acceleration function. With the heating wind mode, SPH particles are created with zero initial velocity with respect to the star, but instead wind particles are given an internal energy based on the integrated mechanical luminosity of the star. This mode is designed to be used on longer timescales and larger spatial scales compared to the other two modes and assumes that the star is embedded in a gas cloud.
    [Show full text]
  • Coronae, Heliospheres and Astrospheres
    Coronae, Heliospheres and Astrospheres Heliophysics Summer School, Boulder CO, 2018 Outline Part I: I. Solar Vs. stellar physics II. Stellar evolution III. Coronae and winds IV. Stellar environments - astrospheres Part II V.Stellar evolution and magnetized winds VI. Stellar mass-loss rates and stellar spin-down VII. Flares VIII.Exoplanets and planet habitability Material mostly based on Volume IV chapters 2,3,4 Part I The Solar-stellar connection SDO observations of the Sun Photometry - measuring the intensity of the light Spectrometry - measuring the intensity of particular wavelength Transmission spectra - E=hv How many photons with certain energy (thus wavelength) we observe SDO AIA, SDO EVE 9-33 nm <30 nm Fermi JWST 0.6 to 27 micrometer Chandra Kepler systems observed as of Jan 2016 Solar Physics: 1. High-resolution global observations 2. High-cadence observations of temporal evolution 3. Multi-wavelength observations 4. In-situ observations of the interplanetary environment 5. Detailed and constrained models 6. Information only about one star Stellar Astrophysics: 1. Statistical information on many stars 2. Data on different spectral types 3. Data on stellar evolution of each type, including solar analogs 4. Information about planetary systems 5. Limited knowledge about specific parameters 6. Limited knowledge about stellar winds and interplanetary environments 7. Unconstrained models Stellar Evolution Stellar Evolution (yes, some figures are from Wikipedia…) Sun Credit: ESO Star-forming regions - molecular clouds Protoplanetary disk What drives stars? Nuclear Fusion Heavy products sink to the core Main Sequence: H—> He Post-main Sequence: Mid-size stars 1. sub-giant phase - H burning in the H shell 2.
    [Show full text]
  • On Possible Types of Magnetospheres of Hot Jupiters
    On possible types of magnetospheres of hot Jupiters Zhilkin, A.G.∗, Bisikalo D.V. Institute of Astronomy of Russian Academy of Sciences, Moscow, Russia ABSTRACT. We show that the orbits of exoplanets of the “hot Jupiter” type, as a rule, are located close to the Alf´ven point of the stellar wind of the parent star. At this, many hot Jupiters can be located in the sub-Alf´ven zone in which the magnetic pressure of the stellar wind exceeds its dynamic pressure. Therefore, magnetic field of the wind must play an extremely important role for the flow of the stellar wind around the atmospheres of the hot Jupiters. This factor must be considered both in theoretical models and in the interpretation of observational data. The analysis shows that many typical hot Jupiters should have shock-less intrinsic magnetospheres, which, apparently, do not have counterparts in the Solar System. Such magnetospheres are characterized, primarily, by the absence of the bow shock, while the magnetic barrier (ionopause) is formed by the induced currents in the upper layers of the ionosphere. We confirmed this inference by the three-dimensional numerical simulation of the flow of the parent star stellar wind around the hot Jupiter HD 209458b in which we took into account both proper magnetic field of the planet and magnetic field of the wind. 1 Introduction When celestial bodies that have their own magnetic field interact with surrounding ionized matter, they create a cavity around themselves, which is named magnetosphere. In particular, such magnetospheres have Solar System planets blown over by the solar wind plasma [1].
    [Show full text]
  • Predicting Radio Emission from the Newborn Hot Jupiter V830 Tau B and Its Host Star A
    Astronomy & Astrophysics manuscript no. radio_V830Tau_vidotto c ESO 2021 July 7, 2021 Predicting radio emission from the newborn hot Jupiter V830 Tau b and its host star A. A. Vidotto1 and J.-F. Donati2; 3 1 School of Physics, Trinity College Dublin, University of Dublin, Ireland 2 Université de Toulouse, UPS-OMP, IRAP, 14 avenue E. Belin, Toulouse F-31400, France 3 CNRS, IRAP / UMR 5277, 14 avenue E. Belin, Toulouse F-31400, France Received date / Accepted date ABSTRACT Magnetised exoplanets are expected to emit at radio frequencies analogously to the radio auroral emission of Earth and Jupiter. Here, we predict the radio emission from V830 Tau b, the youngest (2 Myr) detected exoplanet to date. We model the wind of its host star using three-dimensional magnetohydrodynamics simulations that take into account the reconstructed stellar surface magnetic field. Our simulations allow us to constrain the local conditions of the environment surrounding V830 Tau b that we use to then compute its radio emission. We estimate average radio flux densities of 6 to 24 mJy, depending on the assumption of the radius of the planet (one or two Jupiter radii). These radio fluxes are not constant along one planetary orbit, and present peaks that are up to twice the average values. We show here that these fluxes are weakly dependent (a factor of 1.8) on the assumed polar planetary magnetic field (10 to 100 G), opposed to the maximum frequency of the emission, which ranges from 18 to 240 MHz. We also estimate the thermal radio emission from the stellar wind.
    [Show full text]
  • William Pendry Bidelman (1918-2011)
    William Pendry Bidelman (1918–2011)1 Howard E. Bond2 Received ; accepted arXiv:1609.09109v1 [astro-ph.SR] 28 Sep 2016 1Material for this article was contributed by several family members, colleagues, and former students, including: Billie Bidelman Little, Joseph Little, James Caplinger, D. Jack MacConnell, Wayne Osborn, George W. Preston, Nancy G. Roman, and Nolan Walborn. Any opinions stated are those of the author. 2Department of Astronomy & Astrophysics, Pennsylvania State University, University Park, PA 16802; [email protected] –2– ABSTRACT William P. Bidelman—Editor of these Publications from 1956 to 1961—passed away on 2011 May 3, at the age of 92. He was one of the last of the masters of visual stellar spectral classification and the identification of peculiar stars. I re- view his contributions to these subjects, including the discoveries of barium stars, hydrogen-deficient stars, high-galactic-latitude supergiants, stars with anomalous carbon content, and exotic chemical abundances in peculiar A and B stars. Bidel- man was legendary for his encyclopedic knowledge of the stellar literature. He had a profound and inspirational influence on many colleagues and students. Some of the bizarre stellar phenomena he discovered remain unexplained to the present day. Subject headings: obituaries (W. P. Bidelman) –3– William Pendry Bidelman—famous among his astronomical colleagues and students for his encyclopedic knowledge of stellar spectra and their peculiarities—passed away at the age of 92 on 2011 May 3, in Murfreesboro, Tennessee. He was Editor of these Publications from 1956 to 1961. Bidelman was born in Los Angeles on 1918 September 25, but when the family fell onto hard financial times, his mother moved with him to Grand Forks, North Dakota in 1922.
    [Show full text]
  • Mass Loss from Hot Massive Stars
    Astronomy and Astrophysics Review manuscript No. (will be inserted by the editor) Joachim Puls Jorick S. Vink Francisco Najarro· · Mass loss from hot massive stars Received: date Abstract Mass loss is a key process in the evolution of massive stars, and must be understood quantitatively if it is to be successfully included in broader as- trophysical applications such as galactic and cosmic evolution and ionization. In this review, we discuss various aspects of radiation driven mass loss, both from the theoretical and the observational side. We focus on developments in the past decade, concentrating on the winds from OB-stars, with some excur- sions to the winds from Luminous Blue Variables (including super-Eddington, continuum-driven winds), winds from Wolf-Rayet stars, A-supergiants and Cen- tral Stars of Planetary Nebulae. After recapitulating the 1-D, stationary standard model of line-driven winds, extensions accounting for rotation and magnetic fields are discussed. Stationary wind models are presented that provide theoretical pre- dictions for the mass-loss rates as a function of spectral type, metallicity, and the proximity to the Eddington limit. The relevance of the so-called bi-stability jump is outlined. We summarize diagnostical methods to infer wind properties from observations, and compare the results from corresponding campaigns (in- cluding the VLT-FLAMES survey of massive stars) with theoretical predictions, featuring the mass loss-metallicity dependence. Subsequently, we concentrate on two urgent problems, weak winds and wind-clumping, that have been identified from various diagnostics and that challenge our present understanding of radia- tion driven winds. We discuss the problems of “measuring” mass-loss rates from weak winds and the potential of the NIR Brα -line as a tool to enable a more pre- arXiv:0811.0487v1 [astro-ph] 4 Nov 2008 Joachim Puls Universit¨atssternwarte M¨unchen, Scheinerstr.
    [Show full text]