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Collapsars As a Major Source of R-Process Elements
Collapsars as a major source of r-process elements Daniel M. Siegel1;2;3;4;5, Jennifer Barnes1;2;5 & Brian D. Metzger1;2 1Department of Physics, Columbia University, New York, NY, USA 2Columbia Astrophysics Laboratory, Columbia University, New York, NY, USA 3Perimeter Institute for Theoretical Physics, Waterloo, Ontario, Canada 4Department of Physics, University of Guelph, Guelph, Ontario, Canada 5NASA Einstein Fellow The production of elements by rapid neutron capture (r-process) in neutron-star mergers is expected theoretically and is supported by multimessenger observations1–3 of gravitational- wave event GW170817: this production route is in principle sufficient to account for most of the r-process elements in the Universe4. Analysis of the kilonova that accompanied GW170817 identified5, 6 delayed outflows from a remnant accretion disk formed around the newly born black hole7–10 as the dominant source of heavy r-process material from that event9, 11. Sim- ilar accretion disks are expected to form in collapsars (the supernova-triggering collapse of rapidly rotating massive stars), which have previously been speculated to produce r-process elements12, 13. Recent observations of stars rich in such elements in the dwarf galaxy Retic- arXiv:1810.00098v2 [astro-ph.HE] 14 Aug 2020 ulum II14, as well as the Galactic chemical enrichment of europium relative to iron over longer timescales15, 16, are more consistent with rare supernovae acting at low stellar metal- licities than with neutron-star mergers. Here we report simulations that show that collapsar accretion disks yield sufficient r-process elements to explain observed abundances in the Uni- 1 verse. Although these supernovae are rarer than neutron- star mergers, the larger amount of material ejected per event compensates for the lower rate of occurrence. -
2.5 Dimensional Radiative Transfer Modeling of Proto-Planetary Nebula Dust Shells with 2-DUST
PLANETARY NEBULAE: THEIR EVOLUTION AND ROLE IN THE UNIVERSE fA U Symposium, Vol. 209, 2003 S. Kwok, M. Dopita, and R. Sutherland, eds. 2.5 Dimensional Radiative Transfer Modeling of Proto-Planetary Nebula Dust Shells with 2-DUST Toshiya Ueta & Margaret Meixner Department of Astronomy, MC-221 , University of fllinois at Urbana-Champaign, Urbana, IL 61801, USA ueta@astro. uiuc.edu, [email protected]. edu Abstract. Our observing campaign of proto-planetary nebula (PPN) dust shells has shown that the structure formation in planetary nebu- lae (PNs) begins as early as the late asymptotic. giant branch (AGB) phase due to intrinsically axisymmetric superwind. We describe our latest numerical efforts with a multi-dimensional dust radiative transfer code, 2-DUST, in order to explain the apparent dichotomy of the PPN mor- phology at the mid-IR and optical, which originates most likely from the optical depth of the shell combined with the effect of inclination angle. 1. Axisymmetric PPN Shells and Single Density Distribution Model Since the PPN phase predates the period of final PN shaping due to interactions of the stellar winds, we can investigate the origin of the PN structure formation by observationally establishing the material distribution in the PPN shells, in which pristine fossil records of AGB mass loss histories are preserved in their benign circumstellar environment. Our mid-IR and optical imaging surveys of PPN candidates in the past several years (Meixner at al. 1999; Ueta et al. 2000) have revealed that (1) PPN shells are intrinsically axisymmetric most likely due to equatorially-enhanced superwind mass loss at the end of the AGB phase and (2) varying degrees of equatorial enhancement in superwind would result in different optical depths of the PPN shell, thereby heavily influencing the mid-IR and optical morphologies. -
Spectroscopic Analysis of Accretion/Ejection Signatures in the Herbig Ae/Be Stars HD 261941 and V590 Mon T Moura, S
Spectroscopic analysis of accretion/ejection signatures in the Herbig Ae/Be stars HD 261941 and V590 Mon T Moura, S. Alencar, A. Sousa, E. Alecian, Y. Lebreton To cite this version: T Moura, S. Alencar, A. Sousa, E. Alecian, Y. Lebreton. Spectroscopic analysis of accretion/ejection signatures in the Herbig Ae/Be stars HD 261941 and V590 Mon. Monthly Notices of the Royal Astronomical Society, Oxford University Press (OUP): Policy P - Oxford Open Option A, 2020, 494 (3), pp.3512-3535. 10.1093/mnras/staa695. hal-02523038 HAL Id: hal-02523038 https://hal.archives-ouvertes.fr/hal-02523038 Submitted on 16 May 2020 HAL is a multi-disciplinary open access L’archive ouverte pluridisciplinaire HAL, est archive for the deposit and dissemination of sci- destinée au dépôt et à la diffusion de documents entific research documents, whether they are pub- scientifiques de niveau recherche, publiés ou non, lished or not. The documents may come from émanant des établissements d’enseignement et de teaching and research institutions in France or recherche français ou étrangers, des laboratoires abroad, or from public or private research centers. publics ou privés. MNRAS 000,1–24 (2019) Preprint 27 February 2020 Compiled using MNRAS LATEX style file v3.0 Spectroscopic analysis of accretion/ejection signatures in the Herbig Ae/Be stars HD 261941 and V590 Mon T. Moura1?, S. H. P. Alencar1, A. P. Sousa1;2, E. Alecian2, Y. Lebreton3;4 1Universidade Federal de Minas Gerais, Departamento de Física, Av. Antônio Carlos 6627, 31270-901, Brazil 2Univ. Grenoble Alpes, IPAG, F-38000 Grenoble, France 3LESIA, Observatoire de Paris, PSL Research University, CNRS, Sorbonne Universités, UPMC Univ. -
The Variable Stars and Blue Horizontal Branch of the Metal-Rich Globular Cluster NGC 6441
View metadata, citation and similar papers at core.ac.uk brought to you by CORE provided by CERN Document Server The Variable Stars and Blue Horizontal Branch of the Metal-Rich Globular Cluster NGC 6441 Andrew C. Layden1;2 Physics & Astronomy Dept., Bowling Green State Univ., Bowling Green, OH 43403, U.S.A. Laura A. Ritter Department of Astronomy, University of Michigan, Ann Arbor, MI 48109-1090, U.S.A. Douglas L. Welch1,TracyM.A.Webb1 Department of Physics & Astronomy, McMaster University, Hamilton, Ontario L8S 4M1, Canada ABSTRACT We present time-series VI photometry of the metal-rich ([Fe=H]= 0:53) globular − cluster NGC 6441. Our color-magnitude diagram shows that the extended blue horizontal branch seen in Hubble Space Telescope data exists in the outermost reaches of the cluster. About 17% of the horizontal branch stars lie blueward and brightward of the red clump. The red clump itself slopes nearly parallel to the reddening vector. A component of this slope is due to differential reddening, but part is intrinsic. The blue horizontal branch stars are more centrally concentrated than the red clump stars, suggesting mass segregation and a possible binary origin for the blue horizontal branch stars. We have discovered 50 new variable stars near NGC 6441, among ∼ them eight or more RR Lyrae stars which are highly-probable cluster members. Comprehensive period searches over the range 0.2–1.0 days yielded unusually long periods (0.5–0.9 days) for the fundamental pulsators compared with field RR Lyrae of the same metallicity. Three similar long-period RR Lyrae are known in other metal-rich globulars. -
White Dwarfs
Chandra X-Ray Observatory X-Ray Astronomy Field Guide White Dwarfs White dwarfs are among the dimmest stars in the universe. Even so, they have commanded the attention of astronomers ever since the first white dwarf was observed by optical telescopes in the middle of the 19th century. One reason for this interest is that white dwarfs represent an intriguing state of matter; another reason is that most stars, including our sun, will become white dwarfs when they reach their final, burnt-out collapsed state. A star experiences an energy crisis and its core collapses when the star's basic, non-renewable energy source - hydrogen - is used up. A shell of hydrogen on the edge of the collapsed core will be compressed and heated. The nuclear fusion of the hydrogen in the shell will produce a new surge of power that will cause the outer layers of the star to expand until it has a diameter a hundred times its present value. This is called the "red giant" phase of a star's existence. A hundred million years after the red giant phase all of the star's available energy resources will be used up. The exhausted red giant will puff off its outer layer leaving behind a hot core. This hot core is called a Wolf-Rayet type star after the astronomers who first identified these objects. This star has a surface temperature of about 50,000 degrees Celsius and is A composite furiously boiling off its outer layers in a "fast" wind traveling 6 million image of the kilometers per hour. -
Planetary Nebulae Jacob Arnold AY230, Fall 2008
Jacob Arnold Planetary Nebulae Jacob Arnold AY230, Fall 2008 1 PNe Formation Low mass stars (less than 8 M) will travel through the asymptotic giant branch (AGB) of the familiar HR-diagram. During this stage of evolution, energy generation is primarily relegated to a shell of helium just outside of the carbon-oxygen core. This thin shell of fusing He cannot expand against the outer layer of the star, and rapidly heats up while also quickly exhausting its reserves and transferring its head outwards. When the He is depleted, Hydrogen burning begins in a shell just a little farther out. Over time, helium builds up again, and very abruptly begins burning, leading to a shell-helium-flash (thermal pulse). During the thermal-pulse AGB phase, this process repeats itself, leading to mass loss at the extended outer envelope of the star. The pulsations extend the outer layers of the star, causing the temperature to drop below the condensation temperature for grain formation (Zijlstra 2006). Grains are driven off the star by radiation pressure, bringing gas with them through collisions. The mass loss from pulsating AGB stars is oftentimes referred to as a wind. For AGB stars, the surface gravity of the star is quite low, and wind speeds of ~10 km/s are more than sufficient to drive off mass. At some point, a super wind develops that removes the envelope entirely, a phenomenon not yet fully understood (Bernard-Salas 2003). The central, primarily carbon-oxygen core is thus exposed. These cores can have temperatures in the hundreds of thousands of Kelvin, leading to a very strong ionizing source. -
Plotting Variable Stars on the H-R Diagram Activity
Pulsating Variable Stars and the Hertzsprung-Russell Diagram The Hertzsprung-Russell (H-R) Diagram: The H-R diagram is an important astronomical tool for understanding how stars evolve over time. Stellar evolution can not be studied by observing individual stars as most changes occur over millions and billions of years. Astrophysicists observe numerous stars at various stages in their evolutionary history to determine their changing properties and probable evolutionary tracks across the H-R diagram. The H-R diagram is a scatter graph of stars. When the absolute magnitude (MV) – intrinsic brightness – of stars is plotted against their surface temperature (stellar classification) the stars are not randomly distributed on the graph but are mostly restricted to a few well-defined regions. The stars within the same regions share a common set of characteristics. As the physical characteristics of a star change over its evolutionary history, its position on the H-R diagram The H-R Diagram changes also – so the H-R diagram can also be thought of as a graphical plot of stellar evolution. From the location of a star on the diagram, its luminosity, spectral type, color, temperature, mass, age, chemical composition and evolutionary history are known. Most stars are classified by surface temperature (spectral type) from hottest to coolest as follows: O B A F G K M. These categories are further subdivided into subclasses from hottest (0) to coolest (9). The hottest B stars are B0 and the coolest are B9, followed by spectral type A0. Each major spectral classification is characterized by its own unique spectra. -
• Classifying Stars: HR Diagram • Luminosity, Radius, and Temperature • “Vogt-Russell” Theorem • Main Sequence • Evolution on the HR Diagram
Stars • Classifying stars: HR diagram • Luminosity, radius, and temperature • “Vogt-Russell” theorem • Main sequence • Evolution on the HR diagram Classifying stars • We now have two properties of stars that we can measure: – Luminosity – Color/surface temperature • Using these two characteristics has proved extraordinarily effective in understanding the properties of stars – the Hertzsprung- Russell (HR) diagram If we plot lots of stars on the HR diagram, they fall into groups These groups indicate types of stars, or stages in the evolution of stars Luminosity classes • Class Ia,b : Supergiant • Class II: Bright giant • Class III: Giant • Class IV: Sub-giant • Class V: Dwarf The Sun is a G2 V star Luminosity versus radius and temperature A B R = R R = 2 RSun Sun T = T T = TSun Sun Which star is more luminous? Luminosity versus radius and temperature A B R = R R = 2 RSun Sun T = T T = TSun Sun • Each cm2 of each surface emits the same amount of radiation. • The larger stars emits more radiation because it has a larger surface. It emits 4 times as much radiation. Luminosity versus radius and temperature A1 B R = RSun R = RSun T = TSun T = 2TSun Which star is more luminous? The hotter star is more luminous. Luminosity varies as T4 (Stefan-Boltzmann Law) Luminosity Law 2 4 LA = RATA 2 4 LB RBTB 1 2 If star A is 2 times as hot as star B, and the same radius, then it will be 24 = 16 times as luminous. From a star's luminosity and temperature, we can calculate the radius. -
SHELL BURNING STARS: Red Giants and Red Supergiants
SHELL BURNING STARS: Red Giants and Red Supergiants There is a large variety of stellar models which have a distinct core – envelope structure. While any main sequence star, or any white dwarf, may be well approximated with a single polytropic model, the stars with the core – envelope structure may be approximated with a composite polytrope: one for the core, another for the envelope, with a very large difference in the “K” constants between the two. This is a consequence of a very large difference in the specific entropies between the core and the envelope. The original reason for the difference is due to a jump in chemical composition. For example, the core may have no hydrogen, and mostly helium, while the envelope may be hydrogen rich. As a result, there is a nuclear burning shell at the bottom of the envelope; hydrogen burning shell in our example. The heat generated in the shell is diffusing out with radiation, and keeps the entropy very high throughout the envelope. The core – envelope structure is most pronounced when the core is degenerate, and its specific entropy near zero. It is supported against its own gravity with the non-thermal pressure of degenerate electron gas, while all stellar luminosity, and all entropy for the envelope, are provided by the shell source. A common property of stars with well developed core – envelope structure is not only a very large jump in specific entropy but also a very large difference in pressure between the center, Pc, the shell, Psh, and the photosphere, Pph. Of course, the two characteristics are closely related to each other. -
(GK 1; Pr 1.1, 1.2; Po 1) What Is a Star?
1) Observed properties of stars. (GK 1; Pr 1.1, 1.2; Po 1) What is a star? Why study stars? The sun Age of the sun Nearby stars and distribution in the galaxy Populations of stars Clusters of stars Distances and characterization of stars Luminosity and flux Magnitudes Parallax Standard candles Cepheid variables SN Ia 2) The HR diagram and stellar masses (GK 2; Pr 1.4; Po 1) Colors of stars B-V Blackbody emission HR Diagram Interpretation of HR diagram stellar radii kinds of stars red giants white dwarfs planetary nebulae AGB stars Cepheids Horizontal branch evolutionary sequence turn off mass and ages Masses from binaries Circular orbits solution General solution Spectroscopic binaries Eclipsing spectroscopic binaries Empirical mass luminosity relation 3) Spectroscopy and abundances (GK 1; Pr 2) Stellar spectra OBFGKM Atomic physics H atom others Spectral types Temperature and spectra Boltzmann equation for levels Saha equation for ionization ionization stages Rotation Stellar Abundances More about ionization stages, e.g. Ca and H Meteorite abundances Standard solar set Abundances in other stars and metallicity 4) Hydrostatic balance, Virial theorem, and time scales (GK 3,4; Pr 1.3, 2; Po 2,8) Assumptions - most of the time Fully ionized gas except very near surface where partially ionized Spherical symmetry Broken by e.g., convection, rotation, magnetic fields, explosion, instabilities, etc Makes equations a lot easier Limits on rotation and magnetic fields Homogeneous composition at birth Isolation (drop this later in course) Thermal -
The Hr Diagram for Late-Type Nearby Stars
379 THE H-R DIAGRAM FOR LATE-TYPE NEARBY STARS AS A FUNCTION OF HELIUM CONTENT AND METALLICITY 1 2 3 2 1 1 Y. Lebreton , M.-N. Perrin ,J.Fernandes ,R.Cayrel ,G.Cayrel de Strob el , A. Baglin 1 Observatoire de Paris, Place J. Janssen - 92195 Meudon Cedex, France 2 Observatoire de Paris, 61 Avenue de l'Observatoire - 75014 Paris, France 3 Observat orio Astron omico da Universidade de Coimbra, 3040 Coimbra, Portugal Key words: Galaxy: solar neighb ourho o d; stars: ABSTRACT abundances; stars: low-mass; stars: HR diagram; Galaxy: abundances. Recent theoretical stellar mo dels are used to discuss the helium abundance of a numberoflow-mass stars for which the p osition in the Hertzsprung-Russell di- 1. INTRODUCTION agram and the metallicity are known with high accu- racy. The knowledge of the initial helium abundance of Hipparcos has provided very high quality parallaxes stars b orn in di erent sites with di erent metallicities of a sample of a hundred disk stars, of typeFtoK,lo- is of great imp ortance for many astrophysical stud- cated in the solar neighb ourho o d. Among these stars ies. The lifetime of a star and its internal structure we have carefully selected those for which detailed very much dep end on its initial helium content and sp ectroscopic analysis has provided e ective temp er- this has imp ortant consequences not only for stellar ature and [Fe/H] ratio with a high accuracy. astrophysics but also in cosmology or in studies of the chemical evolution of galaxies. Wehave calculated evolved stellar mo dels and their Direct measurement of the helium abundance in the asso ciated iso chrones in a large range of mass, for photosphere of a low mass star cannot b e made since several values of the metallicity and of the helium there are no helium lines in the sp ectra. -
Spring Semester 2010 Exam 3 – Answers a Planetary Nebula Is A
ASTR 1120H – Spring Semester 2010 Exam 3 – Answers PART I (70 points total) – 14 short answer questions (5 points each): 1. What is a planetary nebula? How long can a planetary nebula be observed? A planetary nebula is a luminous shell of gas ejected from an old, low- mass star. It can be observed for approximately 50,000 years; after that it its gases will have ceased to glow and it will simply fade from view. 2. What is a white dwarf star? How does a white dwarf change over time? A white dwarf star is a low-mass star that has exhausted all its nuclear fuel and contracted to a size roughly equal to the size of the Earth. It is supported against further contraction by electron degeneracy pressure. As a white dwarf ages, though, it will radiate, thereby cooling and getting less luminous. 3. Name at least three differences between Type Ia supernovae and Type II supernovae. Type Ia SN are intrinsically brighter than Type II SN. Type II SN show hydrogen emission lines; Type Ia SN don't. Type II SN are thought to arise from core collapse of a single, massive star. Type Ia SN are thought to arise from accretion from a binary companion onto a white dwarf star, pushing the white dwarf beyond the Chandrasekhar limit and causing a nuclear detonation that destroys the white dwarf. 4. Why was SN 1987A so important? SN 1987A was the first naked-eye SN in almost 400 years. It was the first SN for which a precursor star could be identified.