“KINEMATICS of the ORION NEBULA CLUSTER: VELOCITY SUBSTRUCTURE and SPECTROSCOPIC BINARIES” (2009, Apj, 697, 1103)∗

ERRATUM: “KINEMATICS OF THE ORION NEBULA CLUSTER: VELOCITY SUBSTRUCTURE AND SPECTROSCOPIC BINARIES” (2009, ApJ, 697, 1103)∗

John J. Tobin1,4, Lee Hartmann1, Gabor Furesz2, Mario Mateo1, and S. Tom Megeath3 1 Department of Astronomy, University of Michigan, Ann Arbor, MI 48109, USA 2 Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA 3 Department Astronomy, University of Toledo, 2801 West Bancroft Street, Toledo, OH 43606, USA Received 2013 July 1; published 2013 July 29

In the published version of this paper, a data problem resulted in the publication of 18 erroneous spectroscopic binaries in Table 11. Table 12 did not have erroneous data published, but the now corrected Table 11 includes seven members of the published version of Table 12. This is because they are detected as spectroscopic binaries from both radial velocity variation and double-peaked correlation functions and we therefore present an updated Table 12. We list the erroneously reported single-line spectroscopic binaries in Table 14. Note that none of the analysis or ﬁgures published in the original version were affected by this issue, the problem occurred in the export of the table data to LaTex format.

∗ This paper includes data gathered with the 6.5 m Magellan Telescopes located at Las Campanas Observatory, Chile; Observations reported here were obtained at the MMT Observatory, a joint facility of the Smithsonian Institution and the University of Arizona. 4 Current address: National Radio Astronomy Observatory, Charlottesville, VA 22903, USA.

1 The Astrophysical Journal, 773:81 (5pp), 2013 August 10 Erratum: 2009,ApJ,697, 1103 d Field ID c [3.6] SB − K K I − V V b N_obs 2 r χ Table 11 ) P 4.09 9.42 3 14.38 1.63 10.55 0.14 1 F1-E1, F1-E2, F11 5.82 23.15 2 18.45 3.16 12.25 0.54 1 F3-E1, F22 5.78 13.32 3 0 0 12.12 0.31 2 F1-E1, F1-E2, F11 5.69 6.86 6 11.52 1.34 7.43 0.74 2 B-1, B-2, F2-E1, F2-E2, S1 9.20 38.23 2 18.04 3.7 11.00 0.82 1 S1 4.23 9.73 35.90 17.91 7.08 3.09 12.23 6 0.28 16.53 2.45 1 11.26 B-1, B-2, C-1, C-2, C-3, F3-E1, F22 0.29 1 C-1, C-2, C-3, F2-E1, F2-E2, F21 9.92 41.46 2 16.35 2.49 11.43 0.29 1 A-1, A-2 7.77 31.83 2 0 0 12.43 0.14 1 F3-E1, F21 6.46 14.87 3 0 0 11.91 0.16 1 F1-E1, F1-E2, F11 8.88 36.80 2 16.42 2.8 10.78 0.38 1 F3-E1, F21 4.59 10.57 3 15.11 1.88 10.65 1.58 1 B-1, B-2, F3-E1, S1 5.71 13.16 3 16.4 2.27 11.85 0.14 1 F1-E1, F1-E2, F11 9.28 15.38 4 13.01 1.2 9.95 0.9 1 F1-E1, F1-E2, S2 5.56 12.80 3 16.87 2.77 11.67 0.27 2 C-1, C-2, C-3, F1-E1, F1-E2, F11 8.06 18.57 3 18.58 2.92 10.43 1.01 1 C-1, C-2, C-3, F2-E1, F2-E2, S2 5.13 20.10 2 18.81 3.5 11.57 0.35 1 B-1, B-2, F2-E1, F2-E2, F22 8.60 19.80 3 16.72 2.7 11.57 0.22 1 F1-E1, F1-E2, F11 4.32 16.54 2 17.8 2.86 12.42 0.31 1 B-1, B-2, F3-E1, F21 5.43 21.41 2 17.52 2.66 12.53 0.16 1 B-1, B-2, F3-E1, F21 5.01 19.55 2 17.79 2.97 12.22 0.25 1 F3-E1, F22 23.57 103.45 2 10.1913.45 0.9 21.90 7.30 4 1.28 0 2 0 10.31 F2-E1, F2-E2, S1 0.8 2 F2-E1, F2-E2, S3 48.03 110.59 3 17.95 2.64 12.09 0.15 2 F2-E1, F2-E2, F22 38.34 88.29 3 0 0 12.48 0 1 F1-E1, F1-E2, F11 19.29 44.41 3 17.04 2.63 11.33 0.23 1 F1-E1, F1-E2, F11 74.52 336.90 2 16.93 2.81 10.64 0.22 1 F3-E1, F21 20.24 32.46 4 14.38 1.62 9.80 0.73 1 F1-E1, F1-E2, S2 30.30 69.77 3 15.73 2.79 10.39 0.18 2 C-1, C-2, C-3 10.01 41.85 2 18.19 3.26 12.05 0.31 1 B-1, B-2, F2-E1, F2-E2, F21 10.48 24.13 3 17.76 3.08 12.03 0.34 1 F1-E1, F1-E2, F11 13.68 12.71 7 15.07 2.27 9.43 0.65 1 C-1, C-2, C-3, F1-E1, F1-E2, S1 83.75 80.21 6 11.06 1.63 7.19 0.2 1 C-1, C-2, C-3, F1-E1, F1-E2, S1 21.95 96.06 2 16.38 2.37 11.19 0.17 1 A-1, A-2, F31 19.21 30.85 4 17.86 2.82 11.60 0.22 2 F2-E1, F2-E2, F3-E1, F22 28.36 125.29 2 17.12 2.57 12.24 0.15 1 A-1, A-2, F31 35.80 56.52 4 17.03 2.55 11.07 0.24 1 D-2, F1-E1, F1-E2, F11 − − − − − − − − − − − − − − − − − − − − − 192.61 297.79 4 14.62 1.67 9.40 0.69 1 A-1, A-2, S3 305.22 470.79 4 14.98 2.25 10.12 0.64 2 B-1, B-2, F2-E1, F2-E2, S3 130.50 202.32 4 17.12 3.02 10.87 0.17 1 F2-E1, F2-E2, F11 − − − − − − − − − − − − − − − − log ( − − − ) Spectroscopic Binaries from Velocity Shifts v 1 1.2 4.0 8.4 18.9 6.3 4.5 31.8 13.6 4.0 4.8 4.9 3.3 14.1 3.2 2.9 2.5 5.3 17.1 5.3 13.1 2.3 17.6 5.6 6.6 7.0 2.4 2.1 12.2 15.5 5.4 5.5 19.9 5.5 2.6 13.1 5.2 22.8 6.5 5.9 5.4 Δ − ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± Max. )(kms 1 1.2 1.1 3.9 3.0 8.2 7.4 18.8 15.1 6.1 10.4 4.2 7.9 31.7 37.4 13.5 18.6 3.9 5.5 4.8 3.8 4.7 6.3 2.6 3.9 13.9 11.9 3.1 2.6 2.9 2.8 2.4 2.9 5.2 5.1 17.1 25.3 5.1 5.2 13.0 12.2 2.2 2.4 17.5 23.6 4.4 4.4 6.5 6.0 6.9 6.9 2.0 3.1 1.8 2.0 11.6 16.2 15.4 32.4 5.2 5.6 5.4 3.8 19.8 15.1 5.4 8.8 2.5 2.4 13.1 17.8 5.0 4.8 22.7 22.7 6.4 5.9 5.8 4.7 5.4 8.1 a − ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± RV 19.1 − 05:01:34.23 0534220-050134 22.7 04:49:49.02 0533364-044949 29.7 05:18:44.64 0534499-051844 47.5 05:44:29.76 0534590-054429 22.7 05:25:19.20 0535056-052519 26.3 05:57:56.88 0534475-055756 9.9 05:16:40.62 0533263-051640 68.0 04:24:39.99 0534242-042439 9.3 05:36:32.40 0533454-053632 18.0 05:17:56.53 0534509-051756 22.5 05:05:30.03 0534593-050530 28.9 06:00:11.30 0534492-060011 71.0 04:50:45.60 0534267-045045 32.8 04:55:03.63 0533478-045503 25.0 05:39:24.12 0534517-053924 23.4 05:29:03.48 0535039-052903 23.4 05:27:35.38 0533298-052735 05:26:34.80 0534279-052634 23.4 05:14:15.54 0533545-051415 49.8 04:40:11.64 0534522-044011 66.2 05:09:55.70 0535046-050955 29.8 04:34:39.58 0534286-043439 20.6 05:29:37.60 0534556-052937 25.5 04:32:33.45 0535050-043233 32.0 05:14:39.84 0534292-051439 28.5 04:50:35.16 0534104-045035 24.5 05:06:01.76 0534561-050601 32.8 05:14:50.28 0535052-051450 22.2 05:13:55.20 0534295-051355 26.0 05:24:19.63 0534120-052419 22.7 05:57:47.06 0534330-055747 47.0 05:02:29.47 0534196-050229 40.6 04:34:03.43 0534203-043403 26.2 04:55:28.83 0534390-045528 39.6 05:32:35.14 0534207-053235 25.4 05:56:14.87 0534444-055614 25.8 05:24:48.52 0534209-052448 26.4 05:35:34.76 0534212-053534 22.2 04:47:58.11 0534452-044758 36.4 04:36:07.71 0534455-043607 31.7 − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − − 05:33:36.441 05:34:49.98 05:34:59.05 05:35:05.60 05:34:47.53 R.A.(J2000)05:33:26.396 † (J2000) Decl. 2massID (km s † † † 05:34:24.246 05:33:45.47 05:34:50.992 05:34:59.322 05:34:49.27 05:34:26.741 05:33:47.802 05:34:51.754 05:35:03.91 05:33:29.832 05:34:27.91 05:33:54.573 05:34:52.21 05:35:04.63 05:34:28.677 05:34:55.603 05:35:05.040 05:34:29.24 05:34:10.45 05:34:56.136 05:35:05.21 05:34:29.50 05:34:12.029 05:34:33.012 05:34:19.674 05:34:20.315 05:34:39.039 05:34:20.730 05:34:44.447 05:34:20.990 05:34:21.237 05:34:45.244 05:34:22.078 05:34:45.549

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Classification of initial state granularity via 2D Fourier expansion

This article has been downloaded from IOPscience. Please scroll down to see the full text article. 2013 J. Phys. G: Nucl. Part. Phys. 40 095103 (http://iopscience.iop.org/0954-3899/40/9/095103)

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Please note that terms and conditions apply. IOP PUBLISHING JOURNAL OF PHYSICS G: NUCLEAR AND PARTICLE PHYSICS J. Phys. G: Nucl. Part. Phys. 40 (2013) 095103 (9pp) doi:10.1088/0954-3899/40/9/095103

Classiﬁcation of initial state granularity via 2D Fourier expansion

C E Coleman-Smith1, H Petersen1,2 and R L Wolpert3

1 Department of Physics, Duke University, Durham, NC 27708-0305, USA 2 Frankfurt Institute for Advanced Studies, D-60438 Frankfurt am Main, Germany 3 Department of Statistical Science, Duke University, Durham, NC 27708-0251, USA

E-mail: [email protected]

Received 8 July 2013 Published 29 July 2013 Online at stacks.iop.org/JPhysG/40/095103

Abstract A new method for quantifying fluctuations in the initial state of heavy ion collisions is presented. The initial state energy distribution is decomposed with a set of orthogonal basis functions which include both angular and radial variation. The resulting two-dimensional Fourier coefficients provide additional information about the nature of the initial state fluctuations compared to a purely angular decomposition. We apply this method to ensembles of initial states generated by both Glauber and color glass condensate Monte-Carlo codes. In addition initial state configurations with varying amounts of fluctuations generated by a dynamic transport approach are analyzed to test the sensitivity of the procedure. The results allow for a full characterization of the initial state structures that is useful to discriminate the different initial state models currently in use.

Communicated by Steffen Bass (Some ﬁgures may appear in colour only in the online journal)

Ultra-relativistic nearly-ideal fluid dynamics has proven to be a very successful tool for modeling the bulk dynamics of the hot dense matter formed during a heavy ion collision [17, 19, 20, 28, 35, 38]. The major uncertainty in determining transport properties of the QGP, such as the ratio of shear viscosity to entropy, lies in the specification of the initial conditions of the collision. The initial conditions have been mainly assumed to be smooth distributions that are parametrized implementations of certain physical assumptions (e.g., Glauber/CGC). Within the last two years the importance of including fluctuations in these distributions has been recognized, leading to a whole new set of experimental observations of higher flow coefficients and their correlations [2, 34, 37, 42]. On the theoretical side there has been a lot of effort to refine the previously schematic models with fluctuation inducing corrections and to employ dynamical descriptions of the early non-equilibrium evolution [3, 11, 32, 36]. Hydrodynamical simulations can take these fluctuations into account by generating an ensemble of runs each with a unique initial condition, so-called event by event simulations.

0954-3899/13/095103+09$33.00 © 2013 IOP Publishing Ltd Printed in the UK & the USA 1 The Astrophysical Journal, 773:82 (24pp), 2013 August 20 doi:10.1088/0004-637X/773/2/82 C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

AN INDEPENDENT MEASUREMENT OF THE INCIDENCE OF Mg ii ABSORBERS ALONG GAMMA-RAY BURST SIGHT LINES: THE END OF THE MYSTERY?

A. Cucchiara1, J. X. Prochaska1,G.Zhu2,B.Menard´ 2,3,14, J. P. U. Fynbo4,D.B.Fox5, H.-W. Chen6, K. L. Cooksey7,S.B.Cenko8, D. Perley9, J. S. Bloom8, E. Berger10,N.R.Tanvir11,V.D’Elia12, S. Lopez13, R. Chornock10, and T. de Jaeger13 1 Department of Astronomy and Astrophysics, UCO/Lick Observatory, University of California, 1156 High Street, Santa Cruz, CA 95064, USA; [email protected] 2 Department of Physics and Astronomy, Johns Hopkins University, Baltimore, MD 21218, USA 3 Kavli Institute for the Physics and Mathematics of the Universe, Tokyo University, Kashiwa 277-8583, Japan 4 Dark Cosmology Centre, Niels Bohr Institute, University of Copenhagen, Juliane Maries Vej 30, DK-2100 Copenhagen, Denmark 5 Department of Astronomy and Astrophysics, Pennsylvania State University, University Park, PA 16802, USA 6 Department of Astronomy and Astrophysics, Kavli Institute for Cosmological Physics, University of Chicago, Chicago, IL 60637, USA 7 MIT Kavli Institute for Astrophysics and Space Research, 77 Massachusetts Avenue, 37-685, Cambridge, MA 02139, USA 8 Department of Astronomy, University of California, Berkeley, CA 94720-3411, USA 9 Department of Astronomy, California Institute of Technology, MC 249-17, 1200 East California Blvd., Pasadena, CA 91125, USA 10 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA 11 Department of Physics and Astronomy, University of Leicester, University Road, Leicester LE1 7RH, UK 12 Istituto Nazionale di Astroﬁsica-Osservatorio Astronomico di Roma, Via di Frascati 33, I-00040 Monte Porzio Catone (RM), Italy 13 Departamento de Astronom´ıa, Universidad de Chile, Casilla 36-D, Santiago, Chile Received 2012 November 27; accepted 2013 May 25; published 2013 July 29

ABSTRACT In 2006, Prochter et al. reported a statistically significant enhancement of very strong Mg ii absorption systems intervening the sight lines to gamma-ray bursts (GRBs) relative to the incidence of such absorption along quasar sight lines. This counterintuitive result has inspired a diverse set of astrophysical explanations (e.g., dust, gravitational lensing) but none of these has obviously resolved the puzzle. Using the largest set of GRB afterglow spectra available, we reexamine the purported enhancement. In an independent sample of GRB spectra with a survey path three times larger than Prochter et al., we measure the incidence per unit redshift of 1 Å rest-frame equivalent width Mg ii absorbers at z ≈ 1tobe(z) = 0.18 ± 0.06. This is fully consistent with current estimates for the incidence of such absorbers along quasar sight lines. Therefore, we do not confirm the original enhancement and suggest those results suffered from a statistical fluke. Signatures of the original result do remain in our full sample ((z) shows an ≈1.5 enhancement over (z)QSO), but the statistical significance now lies at ≈90% c.l. Restricting our analysis to the subset of high-resolution spectra of GRB afterglows (which overlaps substantially with Prochter et al.), we still reproduce a statistically significant enhancement of Mg ii absorption. The reason for this excess, if real, is still unclear since there is no connection between the rapid afterglow follow-up process with echelle (or echellette) spectrographs and the detectability of strong Mg ii doublets. Only a larger sample of such high-resolution data will shed some light on this matter. Key words: gamma-ray burst: general – quasars: absorption lines – techniques: spectroscopic Online-only material: color figures

1. INTRODUCTION For simplicity, throughout the paper we will refer to this last category as “strong,” since it is the only one pertinent to this In the last decade, the study of the intergalactic medium work. Mg ii doublet lines have been surveyed extensively from and circumgalactic medium (CGM) has received a great boost z ≈ 0.1–2.5 in the optical passband and now to z = 5.2 with thanks to large spectroscopic surveys of distant quasars, in near-IR spectroscopy (e.g., Steidel & Sargent 1992; Nestor et al. particular the data set provided by the Sloan Digital Sky Survey 2005; Prochter et al. 2006a; Quider et al. 2011; Simcoe et al. (SDSS; York et al. 2000). These objects randomly sample 2011; Zhu & Menard 2012). The results indicate that while the thousands of lines of sight and, being bright background sources weak and strong absorbers incidence show small if any evolu- of light, probe gas, and matter located in foreground objects. tion with redshift, the very strong Mg ii absorbers present an One of the most commonly surveyed set of transitions in increasing trend up to z ∼ 3 before declining at higher redshift quasar spectra is the Mg ii doublet at 2796 and 2803 Å. Its com- (Prochter et al. 2006a; Matejek & Simcoe 2012). This evolution mon detection stems from the large rest wavelength (which rather closely tracks the cosmic star formation history (Prochter makes them easily detectable by most optical spectrographs et al. 2006a; Zhu & Menard 2012), suggesting that some systems when the absorber is located at redshift z = 0.5–2.2), the rela- may be causally connected to ongoing star formation (Menard´ tively high abundance of Mg, and the strength of this resonance- et al. 2011; Matejek & Simcoe 2012), although accurate analy- line doublet. The Mg ii systems are frequently classiﬁed in terms sis of the SDSS survey needs to be carefully taken into account of the rest-frame equivalent width, Wr, of the bluer component in order to avoid technical biases (Lopez´ & Chen 2012). as “weak” (W2796 < 0.3 Å), “strong” (W2796 > 0.3Å)asin For several decades now, strong Mg ii absorption has been Steidel & Sargent (1992) and Churchill et al. (1999), and “very associated with gas in and around galaxies. Early work identiﬁed ∗ strong” (W2796 > 1.0Å,likeinRodr´ıguez Hidalgo et al. 2012). a small sample of L ≈ L galaxies at modest impact parameters (ρ ≈ 10–50 kpc) to quasars exhibiting strong Mg ii absorption 14 Alfred P. Sloan fellow. (Bergeron 1986; Lanzetta et al. 1987; Steidel 1993), although no

1 The Astrophysical Journal, 773:82 (24pp), 2013 August 20 Cucchiara et al.

Table 1 List of Objects Considered for the Mg ii Analysis

a GRB zGRB Telescope Instrument Resolution S/N Reference (Å) 111229A 1.380 Gemini GMOS 5.86.7 This work 111107A 2.893 Gemini GMOS 5.83.5 This work 111008A 4.989 Gemini GMOS 5.83.8 This work 110918A 0.982 Gemini GMOS 5.8 20 This work 110731A 2.830 Gemini GMOS 5.8 26 This work 110726A 1.036 Gemini GMOS 5.89Thiswork 110213B 1.083 Gemini GMOS 3.45Thiswork 110213A 1.460 Bok FAST 6 20 (4) 110205A 2.214 Lick KAST 11 14 (4) 100906A 1.727 Gemini GMOS 5.8 21 This work 100901A 1.408 Gemini GMOS 3.46Thiswork 100814A 1.438 MAGELLAN MagE 1.8 10 This work 100513A 4.798 Gemini GMOS 5.8 17 This work 100418A 0.624 VLT X-Shooter 0.86/0.72/2∗ 12–38 (19) 100414A 1.368 Gemini GMOS 5.8 11 This work 100302A 4.813 Gemini GMOS 5.83Thiswork 100219A 4.667 Gemini GMOS 1.61.2 This work 091208B 1.063 Gemini GMOS 5.826(2) 091109A 3.076 VLT FORS2 13 3 This work 091029 2.752 Gemini GMOS 5.8 39 This work 091024 1.092 Gemini GMOS 5.855(2) 091020A 1.713 NOT ALFOSC 13 7 This work 090926A 2.106 VLT X-Shooter 1.0 15–30 (6) 090902B 1.822 Gemini GMOS 4 14 This work 090812A 2.454 VLT FORS2 13 15 (18) 090529A 2.625 VLT FORS2 13 4 (18) 090519A 3.851 VLT FORS2 13 3 (18) 090516A 4.109 VLT FORS2 13 24 (18) 090426 2.609 Keck LRIS 5.58Thiswork 090424 0.544 Gemini GMOS 5.8 22 This work 090323 3.567 Gemini GMOS 5.8 17 This work 090313 3.375 Gemini GMOS 5.8 11 This work 081222 2.771 Gemini GMOS 5.8 21 This work 081029 3.847 Gemini GMOS 5.842(2) 081008 1.967 Gemini GMOS 3.435(2) 081007 0.529 Gemini GMOS 5.831(2) 080928 1.690 Gemini/VLT GMOS/FORS2 5.8/13 8/25 (2)/(3) 080916A 0.689 VLT FORS1 13 5 (18) 080913A 6.700 VLT FORS2 13 2.5 (13) 080905B 2.374 VLT FORS1 13 13 (18) 080810 3.350 Keck HIRES 0.18 16 (18) 080805 1.505 VLT FORS2 13 3 (3) 080804 2.205 Gemini GMOS 5.817(2) 080721 2.608 TNG Dolores 8.1 9 (18) 080710 0.845 Gemini GMOS 3.435(2) 080707 1.234 VLT FORS1 13 5 (3) 080607 3.036 Keck LRIS 4 11 (3) 080605A 1.639 VLT FORS2 13 30 (3) 080604 1.416 Gemini GMOS 5.84(2) 080603B 2.686 NOT ALFOSC 13 41 (3) 080603A 1.688 Gemini GMOS 5.838(2) 080520 1.545 VLT FORS2 13 5 (3) 080413B 1.100 Gemini GMOS 3.42(2) 080413A 2.433 Gemini GMOS 4 14 (2)/(9) 080411 1.030 VLT FORS1 13 60 (18) 080330 1.513 NOT ALFOSC 13 18 (3) 080319C 1.949 Gemini GMOS 3.44(2) 080319B 0.937 Gemini/VLT GMOS/UVES 5.8/0.13 45/70 (2)/(9) 080310A 2.4272 VLT UVES 0.13 15 (9) 080210 2.6419 VLT FORS2 13 33 (3) 071122 1.141 Gemini GMOS 5.812(2) 071117 1.334 VLT FORS1 13 4 (3) 071112C 0.823 Gemini GMOS 4 3 (2) 071031 2.692 VLT UVES/FORS2 0.13/13 70/40 (3) 071020 2.145 VLT FORS2 13 6 (3)

THE VARIABLE OPTICAL POLARIZATION AND FERMI OBSERVATIONS OF PMN J0948+0022

Joseph R. Eggen, H. Richard Miller, and Jeremy D. Maune Department of Physics and Astronomy, Georgia State University, Atlanta, GA 30303-3083, USA; [email protected] Received 2013 January 3; accepted 2013 June 3; published 2013 July 29

ABSTRACT We report on observations of the γ-ray and optical photopolarimetric behavior of the radio-loud, narrow-line type-1 Seyfert galaxy PMN J0948+0022 over a 27 month period. As this object has recently been suggested to represent a prototype of an emerging class of blazar-like objects, the observed properties are compared to those of blazars. We extract doubling timescales of roughly 4 hr for the optical and γ-ray bands. The rapid microvariability in the optical/near-IR, signiﬁcant and variable optical polarization, and strong yet rapidly variable γ-ray emission we observe for PMN J0948+0022 are all classical observational characteristics associated with blazars. However, since these observations do not show a clear correlation between the γ-ray and optical behavior, they do not offer conclusive proof that the emissive behavior of PMN J0948+0022 is due to a relativistic jet oriented close to our line of sight. Key words: galaxies: active – galaxies: individual (PMN J0948+0022) – galaxies: photometry – galaxies: Seyfert – polarization Online-only material: color ﬁgures, machine-readable table

1. INTRODUCTION timescales (Abdo et al. 2009; Foschini et al. 2012), and microvariability in the optical (Maune et al. 2013). Ikejiri et al. Recently, several members of a sub-class of active galactic (2011) also observed PMN J0948+0022 to exhibit a very high nuclei (AGNs) have been observed with properties that would degree of linear polarization (18.8%) in the optical (V band), have previously been divided among Seyfert galaxies, broad- when the object was very bright (V = 17.028 ± 0.014). line radio galaxies, and blazars. These objects, the radio-loud To date, no comprehensive, long-term program investigat- narrow-line type-1 Seyferts (RL NLS1), possess the standard ing the optical polarimetric/photometric characteristics of RL identifying properties of narrow-line Seyfert-1 (NLS1; Oster- NLS1 has been reported. This paper provides the results brock & Pogge 1985; Kellermann et al. 1989): strong opti- of such a study for the prototype for this class of objects, cal emission of Fe ii, weak emission from forbidden lines (i.e., PMN J0948+0022. [O iii]/Hβ < 3), and FWHM(Hβ) 2000 km s−1 (Goodrich Throughout this manuscript, Julian Dates (JD) are expressed 1989). However, the property of radio loudness (R 10, where as Modified Julian Dates (MJD). The conversion to MJD is R = f5.0 GHz/f4400 Å) (Kellermann et al. 1989) is markedly rare expressed as MJD = JD − 2.45e6. in galaxies of this type, occurring in <7% of such systems (Komossa et al. 2006). High brightness temperature (1013 K), 2. OBSERVATIONS AND DATA REDUCTIONS radio loudness (Zhou et al. 2003), and strong/rapid variability 2.1. Optical Photopolarimetry Data are, however, properties of blazars. It is the combination of these various observational properties that lead many to now believe All optical polarimetric data used in this study were obtained that RL NLS1 and blazars may both play host to relativistic jets. with the 72 inch Perkins telescope at Lowell Observatory in It is now widely accepted that all varieties of AGNs are Flagstaff, Arizona, using the PRISM instrument, which includes manifestations of the same basic phenomenon—accretion of a polarimeter with a rotating half-wave plate. Data were obtained matter onto a supermassive black hole (SMBH) at the center of during several observing runs between 2011 February and a galaxy. The different classes of AGNs that we then observe 2013 April. The specific dates of each observation are given result, to a large degree, from these objects being oriented in Table 1. The observations consisted of a series of two to differently with respect to our line of sight. Blazars, a class four measurements for the Q and U Stokes parameters per of AGNs characterized by strong and variable emission across polarization observation. Each series consisted of four images, all wavelengths and strong and highly variable polarization in each taken at different instrumental position angles (P.A.s)—0◦, the radio and optical (Blandford & Rees 1978), are believed 45◦,90◦, and 135◦—of the wave plate. to result from the orientation of our line of sight near the Corrections to polarimetric values were obtained from two axis of a relativistic jet of particles being emitted from the sources: in-field comparison stars and seperately observed po- central engine of the source. Until recently, blazers were almost larimetric standards, both polarized and unpolarized (Schmidt exclusively observed to be hosted in elliptical galaxies, with et al. 1992). As the camera has a wide field of view (approxi- very few exceptions (McHardy et al 1994). Since the majority mately 14 × 14), we are able to use field stars for interstellar of NLS1 hosts are spiral galaxies, finding evidence of blazar- polarization corrections by subtracting the average percent po- like behavior in such systems would help to fill in a curiously larization of the brightest field stars. Polarized and unpolarized barren demographic of the blazar population. standard stars are used to calibrate corrections for polariza- PMN J0948+0022 is an object that displays the expected tion P.A. and instrumental polarization (typically less than 1%, properties of an NLS1 galaxy as described above (Zhou et al. Jorstad et al. 2010), respectively. 2003), as well as those of blazars, such as strong and vari- The data were reduced and analyzed using in-house scripts, able emission in the radio through γ-ray energies over long which utilize standard packages in the PyRAF 2.0 suite of

1 The Astrophysical Journal, 773:85 (10pp), 2013 August 20 Eggen, Miller, & Maune

Table 1 Photopolarimetric Observations of J0948+0022 Obtained between 2011 February and 2012 May

MJD R Magnitude P EVPA MJD R Magnitude P EVPA (err) (err) (err) (err) (err) (err) 5599.8 19.18 (0.02) 0.86 (0.50) −32.08 (11.3) 6008.7 19.48 (0.02) 1.92 (0.39) −46.7 (0.3) 5602.8 19.18 (0.02) 2.02 (0.30) −36.9 (19.2) 6039.8 19.07 (0.02) 3.89 (0.61) 28.9 (13.5) 5705.7 18.21 (0.02) 1.35 (1.29) 93.4 (8.6) 6040.7 18.89 (0.03) 3.13 (0.36) 9.0 (1.7) 5706.7 17.88 (0.03) 12.31 (1.21) 22.6 (9.3) 6059.7 19.19 (0.03) 0.90 (0.28) 6.2 (62.7) 5708.7 18.90 (0.02) 4.00 (1.51) 11.7 (42.0) 6251.9 18.94 (0.02) 2.29 (0.69) −44.0 (0.1) 5951.8 19.05 (0.03) 1.66 (0.26) −50.74 (2.6) 6253.0 18.68 (0.02) 2.82 (0.58) 15.2 (3.6) 5982.8 18.678 (0.03) 1.70 (0.48) 19.4 (13.6) 6300.0 18.06 (0.03) 1.89 (1.26) 9.0 (6.5) 5983.8 18.82 (0.03) 5.95 (0.50) −59.3 (2.0) 6301.0 18.86 (0.02) 2.62 (0.65) −64.0 (3.0) 5984.7 18.61 (0.03) 2.45 (0.16) 19.9 (5.3) 6302.0 19.00 (0.04) 1.17 (0.28) 79.7 (27.2) 6007.8 18.71 (0.03) 8.17 (0.74) 58.4 (13.5) 6393.8 18.78 (0.02) 1.20 (0.11) 66.14 (1.4)

Notes. Columns: (1) time of the observation in MJD (JD−2.45e6), (2) optical R-band magnitude and (error), (3) percent polarization of target and (error), and (4) EVPA and (error). reduction tools.1 Bias frames were taken at the beginning our optical data, a 7 arcsec aperture radius was used to perform of every night and combined into a master bias that was differential photometry. subtracted from each image. Flat frames were taken at least 2.3. Fermi-LAT Data once per run, using a featureless screen inside the dome. Each position of the wave plate required its own set of flats, which Gamma-ray data were obtained through the Fermi-LAT would later be combined into one master flat per P.A. for public data server. The Large Area Telescope (LAT), on board application to the appropriate science image(s). Cosmic-ray the Fermi Gamma-ray Space Telescope, is a pair-conversion cleaning was performed on all science images, with the threshold detector sensitive to γ-rays in the 20 MeV to several hundred and flux-ratio parameters set to 35 and 5, respectively. Aperture GeV energy range (Atwood et al. 2009). The instrument has photometry was then performed on the calibrated science frames worked almost continuously in all-sky-survey mode since its on an object-by-object basis. An aperture radius of 7 arcsec launch in 2008 June, which allows coverage of the entire was used on all images both to maximize the signal to noise γ-ray sky approximately every 3 hr. The data were reduced and and to maintain consistency with the optical photometry being analyzed using ScienceTools v9r27p1 and instrument response performed on this target by Maune et al. (2013). Use of the functions P7SOURCE_V6. We utilized the likelihood analysis in-field comparison stars compiled by Maune et al. allowed for procedure as described at the FSSC Web site. Photon fluxes one to obtain simultaneous measures of the absolute R-band were calculated using data from MJD 5562 to 6045 (2011 magnitude, as well as the percent polarization (P) and electric January 1–2013 April 25). vector position angle (EVPA). Our data were downloaded from the Fermi Web site on 2013 For ease of comparison with γ-ray data, which are expressed April 29 and cover a region on the sky 15◦ in radius, centered as photon fluxes in this manuscript, optical data were converted on the location of PMN J0948+0022 (2FGL0948.8+0020 from from magnitudes to units of flux (mJy) using the following the Fermi two-year Point Source Catalog), and in an energy equation: range of 100 MeV–300 GeV. Our γ-ray light curve consists of F = 2941 × 10−0.4∗Mag, (1) 112 equally sized bins, each of which is 637,861 s in length, or one-quarter of the lunar synodic period, as our observing runs at where F is the flux in mJy and Mag is the R-band magnitude. Lowell Observatory were centered around the time of the New Moon. The first bin began on 2011 February 3, while the last bin 2.2. Optical and Infrared Photometry with SMARTS ended on 2013 April 25. Only data corresponding to the source ◦ Much of our optical and all of our near-IR (NIR) data were ob- class (evclass = 2) were utilized, with a 52 cutoff rock-angle of ◦ tained by the 1.3 m telescope at the Cerro Tololo Inter-American the spacecraft, while an additional cut utilizing an angle of 100 Observatory under the Small and Moderate Aperture Research from the zenith was imposed so as to minimize the contamina- Telescope System (SMARTS) program. We obtained simultane- tion due to γ-rays coming from Earth’s upper atmosphere. Since ◦ ous data in the optical R and infrared J bands using ANDICAM, PMN J0948+0022 is within 20 of the ecliptic, and the Sun is a which is a dual-channel instrument that uses a dichroic to si- source of γ-rays comparable to our target (Abdo et al. 2011b), multaneously feed optical and IR CCD imagers, allowing the a final cut was used to exclude exposures that occurred when ◦ acquisition of IR data from 0.4 to 2.2 μm. Our limited data set the Sun was within the 10 Radius of Interest (RoI). Photon of J-band images consisted of four NIR images—one for each fluxes and spectral fits were derived using an unbinned max- corresponding optical R-band image—which were flat-fielded, imum likelihood analysis which was accomplished using the overscan-corrected, bias-subtracted, and co-added using stan- ScienceTool gtlike. dard PyRAF/IRAF packages and scripts. To be consistent with In order to accurately measure the flux and spectral parameters of the source, one needs to account for γ-rays emitted from 1 PyRAF is a product of the Space Telescope Science Institute, which is the background. To this end, two models were used: an isotropic operated by AURA for NASA. background model accounting for extragalactic diffuse emission

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Simple effective interaction: infinite nuclear matter and finite nuclei

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Please note that terms and conditions apply. IOP PUBLISHING JOURNAL OF PHYSICS G: NUCLEAR AND PARTICLE PHYSICS J. Phys. G: Nucl. Part. Phys. 40 (2013) 095105 (30pp) doi:10.1088/0954-3899/40/9/095105

Simple effective interaction: inﬁnite nuclear matter and ﬁnite nuclei

B Behera1,XVinas˜ 2, M Bhuyan1,3, T R Routray1, B K Sharma2 andSKPatra3

1 School of Physics, Sambalpur University, Jyotivihar-768 019, India 2 Department d’Estructura i Constitutuents de Materia, University de Barcelona, Diagonal 645, E-08028, Barcelona, Spain 3 Institute of Physics, Sachivalaya Marg, Bhubaneswar-751 005, India

E-mail: [email protected]

Received 7 May 2013 Published 29 July 2013 Online at stacks.iop.org/JPhysG/40/095105

Abstract The mean field properties and equation of state for asymmetric nuclear matter are studied using a simple effective interaction, which has a single finite- range Gaussian term. The study of finite nuclei with this effective interaction is done by constructing a quasilocal energy density functional for which the single-particle equations take the form of Skryme–Hartree–Fock equations. The predictions of binding energies and charge radii of spherical nuclei are found to be compatible with the results of successful mean field models, as well as with the experimental data. (Some figures may appear in colour only in the online journal)

1. Introduction

The last three decades have seen a regular interest in explaining consistently the properties of nuclear matter (NM), finite nuclei and nuclear reactions (nucleon–nucleon, nucleon–nucleus and nucleus–nucleus) with an effective interaction that has the efficiency to describe the two- body system accurately. In this context, the study of nuclear properties from finite nuclei to highly isospin asymmetric NM in a given model is a promising area of current interest. Relativistic and non-relativistic microscopic models, such as Dirac–Brueckner–Hartree–Fock (DBHF) [1–9], Brueckner–Hartree–Fock (BHF) [10–16] and variational calculations using realistic interaction [17, 18] are considered to be standard references in the regime of NM. The ab initio extensions of these models to finite nuclei are still in a preliminary stage. However, an energy density functional based on the microscopic calculations of [16] that reproduces accurately binding energies and charge radii of finite nuclei has been reported recently [19]. Mean field models using effective interactions basically adopt the strategy of fitting the force parameters simultaneously to finite nuclei data and some NM constraints (binding energy per particle, saturation density, incompressibility, symmetry energy, etc).

0954-3899/13/095105+30$33.00 © 2013 IOP Publishing Ltd Printed in the UK & the USA 1 The Astrophysical Journal, 773:78 (21pp), 2013 August 10 doi:10.1088/0004-637X/773/1/78 C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

SYSTEMATICS OF DYNAMICAL MASS EJECTION, NUCLEOSYNTHESIS, AND RADIOACTIVELY POWERED ELECTROMAGNETIC SIGNALS FROM NEUTRON-STAR MERGERS

A. Bauswein1, S. Goriely2, and H.-T. Janka1 1 Max Planck Institute for Astrophysics, Karl-Schwarzschild-Str. 1, D-85748 Garching, Germany 2 Institut d’Astronomie et d’Astrophysique, Universite´ Libre de Bruxelles, C.P. 226, B-1050 Brussels, Belgium Received 2013 February 26; accepted 2013 June 21; published 2013 July 29

ABSTRACT We investigate systematically the dynamical mass ejection, r-process nucleosynthesis, and properties of electromagnetic counterparts of neutron-star (NS) mergers in dependence on the uncertain properties of the nuclear equation of state (EOS) by employing 40 representative, microphysical high-density EOSs in relativistic, hydrodynamical simulations. The crucial parameter determining the ejecta mass is the radius R1.35 of a 1.35 M NS. NSs with −3 −2 smaller R1.35 (“soft” EOS) eject systematically higher masses. These range from ∼10 M to ∼10 M for −3 −2 1.35–1.35 M binaries and from ∼5 × 10 M to ∼2 × 10 M for 1.2–1.5 M systems (with kinetic energies between ∼5 × 1049 erg and 1051 erg). Correspondingly, the bolometric peak luminosities of the optical transients of symmetric (asymmetric) mergers vary between 3 × 1041 erg s−1 and 14 × 1041 erg s−1 (9 × 1041 erg s−1 and 14.5 × 1041 erg s−1) on timescales between ∼2 hr and ∼12 hr. If these signals with absolute bolometric magnitudes from −15.0 to −16.7 are measured, the tight correlation of their properties with those of the merging NSs might provide valuable constraints on the high-density EOS. The r-process nucleosynthesis exhibits a remarkable robustness independent of the EOS, producing a nearly solar abundance pattern above mass number 130. By the r-process content of the Galaxy and the average production per event the Galactic merger rate is limited to 4 × 10−5 yr−1 (4 × 10−4 yr−1) for a soft (stiff) NS EOS, if NS mergers are the main source of heavy r-nuclei. The production ratio of radioactive 232Th to 238U attains a stable value of 1.64–1.67, which does not exclude NS mergers as potential sources of heavy r-material in the most metal-poor stars. Key words: equation of state – hydrodynamics – nuclear reactions, nucleosynthesis, abundances – stars: abundances – stars: neutron Online-only material: color ﬁgures

1. INTRODUCTION help to constrain the still highly uncertain rate of such events in the local universe. Neutron star (NS) merger events are among the most promis- During the merging of two NSs a small fraction of the ing candidates for the first direct measurement of a gravitational- system mass, typically 0.1%–1%, can become gravitationally wave signal with the upcoming Advanced LIGO and VIRGO in- unbound and can be ejected on the dynamical timescale of terferometric instruments (Acernese et al. 2006;Harry&LIGO milliseconds (Ruffert et al. 1997; Rosswog et al. 1999, 2000; Scientific Collaboration 2010), and they are considered as the Ruffert & Janka 2001; Oechslin et al. 2007; Rosswog et al. likely origin of short gamma-ray bursts and their afterglows as a 2013; Piran et al. 2013; Rosswog 2013; Hotokezaka et al. consequence of ultrarelativistic, collimated outflows (see, e.g., 2013). Because such material is likely to possess a high neutron Soderberg et al. 2006; Nakar 2007; Berger 2011; Kann et al. excess, it has been proposed as a possible site for the creation 2011; Fong et al. 2012). Moreover, they are possible sources of of the heaviest, neutron-rich elements, which are formed by the different kinds of electromagnetic signals in the precursor of the rapid neutron capture process (r-process; Lattimer et al. 1977; merging and in its aftermath as a consequence of magnetohy- Eichler et al. 1989) (similarly, also NS–black hole mergers were drodynamical effects, magnetospheric interactions, relativistic suggested as sources of r-process matter; Lattimer & Schramm matter outflows, or NS crust phenomena (Lipunov & Panchenko 1974, 1976). The radioactive decay of these freshly synthesized 1996; Vietri 1996; Li & Paczynski´ 1998; Hansen & Lyutikov r-process nuclei should heat the ejecta and thus lead to an 2001; Troja et al. 2010; Shibata et al. 2011; Nakar & Piran 2011; optical transient (Li & Paczynski´ 1998; Kulkarni 2005; Metzger Tsang et al. 2012; Kyutoku et al. 2012; Zhang 2013; Gao et al. et al. 2010b; Roberts et al. 2011; Goriely et al. 2011). The 2013;Piro2012;Lai2012; Palenzuela et al. 2013; Metzger & properties of such events depend on the fraction of the material Berger 2012). Thermal emission produced by hot ejecta gas, for that can be converted to radioactive species. Moreover, the example, may cause potentially observable optical transients peak luminosity, the timescale to reach the emission peak, and (Li & Paczynski´ 1998; Kulkarni 2005; Metzger et al. 2010b; the effective temperature at the radiation maximum, as well Metzger & Berger 2012), and the interaction of the ejecta cloud as the radio brightness that accompanies the deceleration of with the circumstellar medium is expected to create radio flares the expelled gas during its coasting in the stellar environment, that might be detectable for periods of years (Nakar & Piran depend sensitively on the ejecta mass and expansion velocity. 2011; Piran et al. 2013; Rosswog et al. 2013). Observations of Detailed hydrodynamical merger models are needed to calculate such signals could help to pinpoint the exact celestial locations these quantities and to determine the nucleosynthesis conditions of NS mergers (thus, e.g., supporting the analysis of data taken in the unbound material. by gravitational-wave detectors), and repeated measurements of Concerning their role as sources of heavy elements, binary signals that can be unambiguously linked to NS mergers would NSs collisions have recently moved into the focus of interest

1 The Astrophysical Journal, 773:78 (21pp), 2013 August 10 Bauswein, Goriely, & Janka because the astrophysical sources of the r-process elements have 2013). This can be understood because of several facts. First, the not been identified yet and core-collapse supernova simulations structure of NSs in GR is considerably more compact than that of continue to be unable to yield the extreme conditions for Newtonian stars. For instance, an NS with a gravitational mass forming the heaviest neutron-rich nuclei (Arcones et al. 2007; of 1.35 M described by the LS220 EOS (Lattimer & Swesty Hoffman et al. 2008; Janka et al. 2008; Roberts et al. 2010; 1991) possesses a circumferential radius of 12.6 km, whereas the Hudepohl¨ et al. 2010; Fischer et al. 2010; Wanajo et al. 2011; corresponding Newtonian star has 14.5 km. Second, GR gravity Arcones & Janka 2011; Arcones & Mart´ınez-Pinedo 2011). is stronger and the merging of two NSs is therefore more violent. (For reviews on r-process nucleosynthesis and an overview of The difference can be expressed in terms of the gravitational potential sites, see, e.g., Arnould et al. 2007; Thielemann et al. binding energy of a nucleon on the surface of the considered 2011; Banerjee et al. 2011; Winteler et al. 2012.) In contrast 1.35 M NSs, which is ∼200 MeV in the GR case compared to the situation for supernovae, investigations with growing to only ∼130 MeV for the Newtonian model. Third, GR forces sophistication have confirmed NS merger ejecta as viable sites merger remnants beyond a mass limit to collapse to black holes for strong r-processing (Freiburghaus et al. 1999; Goriely et al. on a dynamical timescale. Such an effect cannot be tracked by 2005; Arnould et al. 2007; Metzger et al. 2010b; Roberts et al. Newtonian models. These differences are of direct relevance for 2011; Goriely et al. 2011; Korobkin et al. 2012). the collision dynamics and the possibility to unbind matter from However, despite this promising situation a variety of aspects the inner and outer crust regions of the merging NSs. need to be clarified before the question can be answered whether It is the purpose of this paper to explore the influence of the NS mergers are a major source or even the dominant source of high-density EOS on the ejecta properties in a systematic way, heavy r-process elements. On the one hand the merger rate i.e., we will determine ejecta masses and the nucleosynthesis and its evolution during the Galactic history are still subject outcome for a large set of NS matter models, applying them in to considerable uncertainties (see, e.g., Abadie et al. 2010 for relativistic NS merger simulations. Most of these EOSs were a compilation of recent estimates), and it is unclear whether already employed in our previous works (Bauswein & Janka NS mergers can explain the early enrichment of the Galaxy by 2012; Bauswein et al. 2012). They were chosen such that they r-process elements as observed in metal-deficient stars (Argast provide as completely as possible a coverage of the possibilities et al. 2004). On the other hand it remains to be determined for NS properties (expressed by corresponding mass–radius- how much mass is ejected in merger events depending on relations) which are compatible with present observational con- the binary parameters and, in particular, depending on the straints (e.g., the 1.97 M NS discovery of Demorest et al. 2010, incompletely known properties of the equation of state (EOS) and more recently 2.01 M by Antoniadis et al. 2013) and theo- of NS matter. It also needs to be understood which fraction retical understanding (Lattimer & Prakash 2010, 2007; Steiner of the ejecta is robustly converted to r-process material and et al. 2010; Hebeler et al. 2010). In our study we will focus whether the final abundances are always compatible with the on symmetric 1.35–1.35 M systems and will compare them solar element distribution, which agrees amazingly well with with asymmetric 1.2–1.5 M mergers. Because population syn- the r-process abundance pattern in metal-poor stars for atomic thesis models (Belczynski et al. 2008) and pulsar observations numbers Z ∼ 55–90 (see, e.g., Sneden et al. 2008). (Thorsett & Chakrabarty 1999; Zhang et al. 2011) suggest that Newtonian as well as relativistic studies showed that the mass the double NS population is strongly dominated by systems ratio has a significant effect on the amount of matter that can of nearly equal-mass stars of about 1.35 M each, the average become unbound (Janka et al. 1999; Rosswog et al. 1999, 2000, NS merger event can be well represented by a 1.35–1.35 M 2013; Ruffert & Janka 2001; Oechslin et al. 2007; Roberts et al. configuration, and a clarification of the EOS dependence of 2011; Goriely et al. 2011; Piran et al. 2013; Korobkin et al. 2012; ejecta masses, r-process yields, and properties of electromag- Rosswog 2013; Hotokezaka et al. 2013). Such investigations, netic counterparts of NS mergers seems to be more important however, were performed only with a few exemplary models for than a wide variation of binary parameters. Nevertheless, we high-density matter in NSs (Rosswog et al. 2000; Oechslin et al. will also present results of a more extended survey of binary 2007; Goriely et al. 2011; Hotokezaka et al. 2013) or even only mass ratios and total masses for some representative EOSs. with a single NS EOS (Roberts et al. 2011; Piran et al. 2013; In our work we will exclusively concentrate on NS–NS Rosswog et al. 2013; Korobkin et al. 2012; Rosswog 2013), mergers, but the discussed phenomena should play a role also although the importance of the nuclear EOS for a quantitative for NS–black hole coalescence (Janka et al. 1999;Lee2000; assessment of the dynamical mass ejection can be concluded Rosswog et al. 2004, 2013; Rosswog 2005, 2013; Faber et al. from published calculations (e.g., Goriely et al. 2011). These 2006; Foucart et al. 2013; Piran et al. 2013) and eccentric NS calculations, however, also suggest that the nuclear abundance mergers (East et al. 2012; East & Pretorius 2012; Rosswog pattern produced by r-processing in the ejecta may be largely et al. 2013; Rosswog 2013). However, while the existence of insensitive to variations of the conditions in the ejecta. double NS systems is established by observations, progenitors It is important to note that quantitatively reliable information of NS–black hole and eccentric NS mergers have not been on the ejecta masses and their dependence on the binary and observed yet and the rates of such types of events are even more EOS properties require general relativistic (GR) simulations. uncertain than those of coalescing binary NSs. In investigating Newtonian results in the literature (Rosswog et al. 1999; Janka the latter we will only consider the phase of dynamical mass et al. 1999; Ruffert & Janka 2001; Roberts et al. 2011; Korobkin ejection between about the time when the two NSs collide until et al. 2012; Rosswog et al. 2013; Piran et al. 2013; Rosswog a few milliseconds later. During this phase hydrodynamical and 2013) exhibit significant quantitative and qualitative differences tidal forces (shock compression, pressure forces, gravitational compared to relativistic models (Oechslin et al. 2007; Goriely interaction) are responsible for the mass shedding of the merging et al. 2011; Hotokezaka et al. 2013). Newtonian calculations objects. Once the remnant has formed, however, differential tend to overestimate the ejecta masses in general (Rosswog et al. rotation is expected to strongly amplify the magnetic fields 1999, 2013; Janka et al. 1999; Ruffert & Janka 2001; Roberts (e.g., Price & Rosswog 2006; Anderson et al. 2008; Liu et al. et al. 2011; Korobkin et al. 2012; Piran et al. 2013; Rosswog 2008; Giacomazzo et al. 2011) and viscous energy dissipation

THE EINSTEIN@HOME SEARCH FOR RADIO PULSARS AND PSR J2007+2722 DISCOVERY

B. Allen1,2,3, B. Knispel1,2, J. M. Cordes4, J. S. Deneva5,J.W.T.Hessels6,7, D. Anderson8, C. Aulbert1,2, O. Bock1,2, A. Brazier4,9, S. Chatterjee4, P. B. Demorest10, H. B. Eggenstein1,2, H. Fehrmann1,2, E. V. Gotthelf11, D. Hammer3, V. M. Kaspi12, M. Kramer13,A.G.Lyne14, B. Machenschalk1,2, M. A. McLaughlin15, C. Messenger1,2, H. J. Pletsch1,2, S. M. Ransom10, I. H. Stairs16, B. W. Stappers14,N.D.R.Bhat17,18, S. Bogdanov11,F.Camilo5,11, D. J. Champion13, F. Crawford19, G. Desvignes20, P. C. C. Freire13, G. Heald6,F.A.Jenet21, P. Lazarus13,K.J.Lee13, J. van Leeuwen6,7, R. Lynch12, M. A. Papa1,2,3, R. Prix1,2,R.Rosen22, P. Scholz12, X. Siemens3, K. Stovall21, A. Venkataraman5, and W. Zhu16 1 Max-Planck-Institut fur¨ Gravitationsphysik, D-30167 Hannover, Germany; [email protected] 2 Leibniz Universitat¨ Hannover, D-30167 Hannover, Germany 3 Department of Physics, University of Wisconsin–Milwaukee, Milwaukee, WI 53211, USA 4 Department of Astronomy, Cornell University, Ithaca, NY 14853, USA 5 Arecibo Observatory, HC3 Box 53995, Arecibo, PR 00612, USA 6 ASTRON, the Netherlands Institute for Radio Astronomy, Postbus 2, 7990 AA, Dwingeloo, The Netherlands 7 Astronomical Institute “Anton Pannekoek,” University of Amsterdam, Science Park 904, 1098 XH Amsterdam, The Netherlands 8 Space Sciences Laboratory, University of California, Berkeley, CA 94720, USA 9 NAIC, Cornell University, Ithaca, NY 14853, USA 10 NRAO (National Radio Astronomy Observatory), Charlottesville, VA 22903, USA 11 Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA 12 Department of Physics, McGill University, Montreal, QC H3A2T8, Canada 13 Max-Planck-Institut fur¨ Radioastronomie, D-53121 Bonn, Germany 14 Jodrell Bank Centre for Astrophysics, School of Physics and Astronomy, University of Manchester, Manchester, M13 9PL, UK 15 Department of Physics, West Virginia University, Morgantown, WV 26506, USA 16 Department of Physics and Astronomy, University of British Columbia, Vancouver, BC V6T 1Z1, Canada 17 Center for Astrophysics and Supercomputing, Swinburne University, Hawthorn, Victoria 3122, Australia 18 International Centre for Radio Astronomy Research, Curtin University, Bentley, WA 6102, Australia 19 Department of Physics and Astronomy, Franklin and Marshall College, Lancaster, PA 17604-3003, USA 20 Department of Astronomy and Radio Astronomy Laboratory, University of California, Berkeley, CA 94720, USA 21 Center for Gravitational Wave Astronomy, University of Texas, Brownsville, TX 78520, USA 22 NRAO, P.O. Box 2, Green Bank, WV 24944, USA Received 2013 February 28; accepted 2013 June 8; published 2013 July 29

ABSTRACT Einstein@Home aggregates the computer power of hundreds of thousands of volunteers from 193 countries, to search for new neutron stars using data from electromagnetic and gravitational-wave detectors. This paper presents a detailed description of the search for new radio pulsars using Pulsar ALFA survey data from the Arecibo Observatory. The enormous computing power allows this search to cover a new region of parameter space; it can detect pulsars in binary systems with orbital periods as short as 11 minutes. We also describe the first Einstein@Home discovery, the 40.8 Hz isolated pulsar PSR J2007+2722, and provide a full timing model. PSR J2007+2722’s pulse profile is remarkably wide with emission over almost the entire spin period. This neutron star is most likely a disrupted recycled pulsar, about as old as its characteristic spin-down age of 404 Myr. However, there is a small chance that it was born recently, with a low magnetic field. If so, upper limits on the X-ray flux suggest but cannot prove that PSR J2007+2722 is at least ∼100 kyr old. In the future, we expect that the massive computing power provided by volunteers should enable many additional radio pulsar discoveries. Key words: binaries: close – gravitational waves – methods: data analysis – pulsars: general – pulsars: individual (PSR J2007+2722) – surveys Online-only material: color figures

1. INTRODUCTION (GW) detectors (Sathyaprakash & Schutz 2009), from radio telescopes (Lyne & Graham-Smith 1998; Lorimer & Kramer Einstein@Home is an ongoing volunteer distributed com- 2004), and from the Large Area Telescope (LAT; Atwood et al. 2009) gamma-ray detector on board the Fermi Satellite. Because puting project (Anderson et al. 2006), launched in early 2005. 23 More than 330,000 members of the general public have “signed the expected signals are weak, and the source parameters are up” their laptop and desktop computers. When otherwise unknown, the sensitivity of the GW searches (Brady et al. 1998; idle, these computers download observational data from the Brady & Creighton 2000) the radio pulsar searches (Brooke Einstein@Home servers, search the data for weak astrophysical et al. 2007), and the gamma-ray searches (Pletsch & Allen 2009; signals, and return the results of the analysis. The collective Pletsch et al. 2012a, 2012b, 2012c) are limited by the available computing power is on par with the largest supercomputers in computing power. the world. 23 Depending upon the type of search, these unknown parameters might The goal of Einstein@Home is to discover neutron stars, include the sky position, spin frequency, spin-down rate, orbital parameters, using data from an international network of gravitational-wave etc.

1 The Astrophysical Journal, 773:91 (32pp), 2013 August 20 Allen et al. Before 2009, Einstein@Home only searched data from the The paper is structured as follows. Section 2 presents a Laser Interferometer Gravitational-Wave Observatory (LIGO; general description of the Einstein@Home computing project, Abramovici et al. 1992; Barish & Weiss 1999; Abbott et al. including its motivation, its history, and its technical design 2009c). So far these searches have not found any sources, and structure. Section 3 is a brief overview of the PALFA but have set new and more sensitive upper limits on possible survey, including its history, the data taking rates, and data continuous gravitational-wave (CW) emissions (Abbott et al. acquisition system. Section 4 is a detailed technical description 2009a, 2009b;Aasietal.2013). These searches are ongoing, of the Einstein@Home search for radio pulsars, starting from the with increasing sensitivity arising from improved data analysis centralized data preparation, through the distributed processing methods (Pletsch & Allen 2009) and better-quality data (Smith on volunteers’ computers, and centralized post-processing. & LIGO Scientific Collaboration 2009). Section 5 describes the discovery of the first Einstein@Home In 2009, Einstein@Home also began searching radio survey radio pulsar, PSR J2007+2722. Section 6 is about the subsequent data from the 305 m Arecibo telescope in Puerto Rico. This is follow-up investigations and studies, including observations at the world’s largest and most sensitive radio telescope, and has multiple frequencies, and accurate determination of the sky discovered a substantial fraction of all known pulsars. Beginning position through gridding and timing. We also discuss the in 2010 December a similar search using data from Parkes evolutionary origin of PSR J2007+2722. This is followed in Observatory in Australia was also started; the differences from Section 7 by a short discussion and conclusion. the Arecibo search and some results are described in Knispel Unless otherwise stated, all coordinates in this paper are in et al. (2013). the J2000 coordinate system, and c denotes the speed of light. Starting in summer 2011, Einstein@Home also began a search for isolated gamma-ray pulsars in data from the Fermi satellite’s 2. THE EINSTEIN@HOME DISTRIBUTED LAT (Atwood et al. 2009); this will be described in future COMPUTING PROJECT publications. 2.1. Volunteer Distributed Computing The Arecibo data are collected by the Pulsar ALFA (PALFA) Consortium using the Arecibo L-band Feed Array (ALFA24). The basic motivation for volunteer distributed computing is For the pulsar survey, ALFA output is fed into fast, broad-band simple: the aggregate computing power owned by the general spectrometers (see Section 3.2); further down the data analysis public exceeds that of universities, and public and private pipeline (see Section 4.1) this enables compensation for the research laboratories, by two to three orders of magnitude. dispersive propagation of pulses from celestial sources. Scientific research whose progress is limited or constrained by The computing capacity of Einstein@Home is used to search computing can benefit from access to even a small fraction of the spectrometer output for signals from neutron stars in short- these resources. This type of research includes both numerical period orbits around companion stars. This is a poorly explored simulation and Monte-Carlo-type exploration of parameter region of parameter space, where other radio-pulsar search spaces, that make no (direct) use of observational data, and data- pipelines lose much or most of their sensitivity. The detection mining and data-analysis efforts which perform deep searches of these pulsars with standard Fourier methods is hampered by through (potentially very large) observational data sets. Doppler smearing of the pulsed signal caused by binary motion Worldwide, there are more than one billion personal comput- during the survey observation (Johnston & Kulkarni 1991). ers (PCs) which are connected to the Internet. These PCs typ- Previous searches (Anderson et al. 1990; Camilo et al. 2000) ically contain x86-architecture central processor units (CPUs) have utilized “acceleration searches” (Johnston & Kulkarni and substantial disk-based and solid-state storage. Many of these 1991), which correct for the part of the binary motion which systems also contain graphics processor units (GPUs) which can can be modeled as a constant acceleration along the line-of- perform floating point calculations one to two orders of magni- sight. These computationally efficient techniques are effective tude faster than a modern CPU core. when the observation time is short compared to the orbital The raw computational capacity of each of these consumer period. Thus, they are insensitive to the most compact systems computers is similar to that of the systems used as building (Ransom et al. 2002). In contrast, the computing power of blocks for computer clusters or research supercomputers. In Einstein@Home enables a full demodulation to be carried out, fact modern research computers are made possible only by giving substantially increased sensitivity to signals from pulsars the economies of scale of the consumer marketplace, which in compact circular orbits with periods below ∼1hr. ensures that the basic components are inexpensive and widely In 2010 August, Einstein@Home announced its first discov- available. But research machines typically consist of hundreds or ery of a new neutron star (Knispel et al. 2010) which appears thousands of these CPUs; volunteer distributed computing offers to be the fastest-spinning “disrupted recycled pulsar” (DRP) so access to hundreds of thousands or millions of these CPUs. far found (Belczynski et al. 2010). In the same month, Ein- stein@Home also discovered a 48 Hz pulsar in a binary sys- 2.2. Constraints on Suitable Computing Problems tem (Knispel et al. 2011). Further Einstein@Home discoveries Volunteer distributed computing is only a suitable solution in Parkes Multi-Beam Pulsar Survey (PMPS) are described in for some computing and data analysis problems: there are Knispel et al. (2013). As of 2013 January, Einstein@Home has both social and technical constraints. To attract volunteers, the discovered almost 50 radio pulsars. research must resonate with the “person in the street.” It must This paper has two purposes. First, it provides a full de- have clear and understandable goals that appeal to the general scription of the Einstein@Home radio pulsar search and post- public and that excite and maintain interest. Experience shows processing pipeline. Second, it provides a detailed description that at least four areas have these qualities: medical research, and full timing solution for the first Einstein@Home discovery, mathematics, climate/environmental science, and astronomy the 40.8 Hz pulsar PSR J2007+2722 (Knispel et al. 2010). and astrophysics. The technical constraints arise because the computers are only 24 http://www.naic.edu/alfa/ connected by the public Internet. This is very different than

KEPLER PHOTOMETRY OF FOUR RADIO-LOUD ACTIVE GALACTIC NUCLEI IN 2010–2012

Ann E. Wehrle1, Paul J. Wiita2, Stephen C. Unwin3, Paolo Di Lorenzo2, Mitchell Revalski2, Daniel Silano2, and Dan Sprague2 1 Space Science Institute, 4750 Walnut Street, Suite 205, Boulder, CO 80301, USA; [email protected] 2 Department of Physics, The College of New Jersey, P.O. Box 7718, Ewing, NJ 08628, USA 3 Jet Propulsion Laboratory, Mail Stop 321-100, 4800 Oak Grove Drive, Pasadena, CA 91109, USA Received 2013 February 28; accepted 2013 May 15; published 2013 July 29

ABSTRACT We have used Kepler photometry to characterize variability in four radio-loud active galactic nuclei (AGN; three quasars and one object tentatively identified as a Seyfert 1.5 galaxy) on timescales from minutes to months, comparable to the light crossing time of the accretion disk around the central supermassive black hole or the base of the relativistic jet. Kepler’s almost continuous observations provide much better temporal coverage than is possible from ground-based observations. We report the first such data analyzed for quasars. We have constructed power spectral densities using eight Kepler quarters of long-cadence (30-minute) data for three AGN, six quarters for one AGN and two quarters of short-cadence (1-minute) data for all four AGN. On timescales longer than about 0.2–0.6 days, we find red noise with mean power-law slopes ranging from −1.8 to −1.2, consistent with the variability originating in turbulence either behind a shock or within an accretion disk. Each AGN has a range of red noise slopes which vary slightly by month and quarter of observation. No quasi-periodic oscillations of astrophysical origin were detected. We detected flares of several days long when brightness increased by 3%–7% in two objects. No flares on timescales of minutes to hours were detected. Our observations imply that the duty cycle for enhanced activity in these radio-loud AGN is small. These well-sampled AGN light curves provide an impetus to develop more detailed models of turbulence in jets and instabilities in accretion disks. Key words: accretion, accretion disks – black hole physics – galaxies: active – galaxies: Seyfert – quasars: general Online-only material: color figures

1. INTRODUCTION The Kepler mission (Borucki et al. 2010; Koch et al. 2010) is uniquely able to probe the innermost regions of AGN and One of the best ways to probe the extremely small regions produce superior light curves through its ability to monitor un- from which the bulk of the energy in active galactic nuclei interrupted by long gaps; it has only brief 1-day gaps every (AGN) is emitted is through the study of their variability in month for data downlink. This capability allows us to study a different bands. The immense powers, non-thermal spectra, and broad range of variability time scales. The overwhelming ma- rapid variability detected across the electromagnetic spectrum jority of monitored blazars display significant microvariability that characterize the class of blazars, i.e., flat spectrum radio (variations of at least 0.03 mag) on timescales of less than a quasars (FSRQs) and BL Lac objects, can only be understood day (see review by Miller & Noble 1996); however, it is not within the framework of matter flowing inward through accre- known if this behavior is characteristic of all blazars because tion disks (ADs) onto supermassive black holes (SMBHs). The observers tend to select the most variable objects to monitor. bulk of this emission emerges from within several gravitational We selected radio-loud AGN in the Kepler field of view for radii of the SMBH; part of the energy is often channeled outward variability monitoring observations obtained over eight quarters via relativistic jets, producing radio-loud AGN. That emission in 2010–2012. Variability probably probes the characteristics of is magnified by Doppler boosting when the jet is pointed within the ADs when the sources are in quiescent faint states and allows a few degrees to our line of sight. Such small viewing angles us to study the synchrotron jet emission when the sources are in cause ordinary radio galaxies and radio loud quasars to appear highly active states. For both radio-loud and radio-quiet quasars, as blazars (e.g., Urry & Padovani 1995). there is a good probability of observing significant variability Optical variability can originate in relativistic jets or in the (a few tenths of a magnitude) over the course of 2 yr (e.g., Pica AD (e.g., Marscher & Gear 1985; Mangalam & Wiita 1993). As et al. 1988; Hawkins 2002; MacLeod et al. 2012). In a study the originating physical processes are different in each case, the of blazars observed with the Palomar Quest survey, Bauer et al. light curves should look different. When FSRQs are in a high (2009) found that 35% of blazars showed V>0.4 mag of state, synchrotron emission from the relativistic jet overwhelms variation over 3.5 yr. emission from the AD. When FSRQs are in a low state, the AD In addition, Kepler is capable of detecting microvariability can become visible, and be recognizable as the Big Blue Bump over the course of several hours at the level of a few percent (Sun & Malkan 1989). During these low and high states, which for targets brighter than about 17th magnitude. Ground-based can differ by 2–5 mag for the most active blazars (optically optical studies have indicated that low-frequency-peaked blazars violent variables), the dominant light source should determine and core-dominated radio galaxies with high polarizations the variability characteristics of the FSRQs, specifically, the frequently exhibit stronger microvariability, as expected if the power spectral densities (PSDs) and possible existence of quasi- jets point toward us and benefit from Doppler boosting. Ordinary periodic oscillations (QPOs). It is not clear from theoretical radio-loud quasars and radio-quiet quasars also show some models, nor from previous observations, if the low and high microvariability, but it is both less frequent and typically weaker state variability states would have PSDs with different slopes. (e.g., Gopal-Krishna et al. 2003;Carinietal.2007;Ram´ırez et al.

1 The Astrophysical Journal, 773:89 (13pp), 2013 August 20 Wehrle et al. 2009; Goyal et al. 2012). Core-dominated radio-loud quasars, QPOs has been suggested for the high declination BL Lac object, which are believed to have jets pointing close to our line-of- S5 0716+714, for some time (e.g., Quirrenbach et al. 1991). A sight, and therefore expected to show substantial fast variations, wavelet analysis of archival optical spectra of S5 0716+714 exhibit much more microvariability than radio-quiet quasars if indicated the presence of five nights with QPOs present at they also exhibit high optical polarizations (Goyal et al. 2012). 0.99 probability, with central periods ranging between 25 This paper is the first report on Kepler data for FSRQs. and 73 minutes (Gupta et al. 2009). New observations of this Mushotzky et al. (2011) reported Kepler monitoring of four same blazar led to an even stronger indication of an ∼15-minute Seyfert galaxies (z<0.09), each for a duration of 2–4 QPO (Rani et al. 2010). quarters; we note that our targets have larger redshifts and are The simplest QPO model assumes that a single hot spot in the substantially fainter (by 2–3 mag). According to their analysis, inner portion of the AD is responsible, and that radiation from all four Seyferts showed some degree of variability over these this disk is directly detected, e.g., when an FSRQ is in a faint, periods and all exhibited very steep red noise components to quiescent state. Given a measured period, P (in seconds) one their PSDs (−2.6 α −3.3). Carini & Ryle (2012) provided can estimate the SMBH mass M via a more in-depth analysis of the Kepler data for one of those M 3.23 × 104P Seyferts, II Zw 229.015, and were able to combine Kepler data = , (1) 3/2 across three quarters because they had sufficient ground-based M (r + a)(1 + z) measurements to normalize the fluxes on different detectors of where r is the hot spot distance in units of GM/c2, a is the Kepler’s camera during different quarters. This allowed them black hole (BH) spin-parameter, and z is the redshift (Gupta to extend the PSD to lower frequencies where they detected a et al. 2009). The shortest periods are obtained for the innermost ∼ flattening below a frequency corresponding to 44 days; their stable circular orbit, and the strongest radiation emerges near best fit to the data in the regime measured by Mushotzky et al. there. For low redshift objects, a 5 hr period, which would − (2011) led them to a shallower PSD slope of around 2.8. Our be very difficult to detect from ground-based observations, but results on FSRQs show significantly flatter PSDs than either rather easy for Kepler to find, corresponds to SMBH masses of 7 8 group; we discuss this in Section 6. 4.0 × 10 and 2.5 × 10 M, for Schwarzschild and extreme In this paper, we first briefly review in Section 2 the physical Kerr BHs, respectively. In practical terms, we need to observe origins of the variability signatures that disks and jets can be for many months to be sensitive to ADs surrounding billion- expected to exhibit, then describe our target sample selection solar-mass BHs because any large hot spots in ADs are probably in Section 3. In Section 4 we explain the steps needed to re- both rare and short-lived, lasting no more than dozens of orbits duce Kepler data on quasars, as it differs significantly from the (∼2 weeks). standard analysis used for the primary Kepler science goal of exoplanet transit detection. We cover the post-processing anal- 2.3. Emission from Jets ysis, including Palomar Observatory imaging and photometry, For core-dominated FSRQs, where the emission is almost in Section 5. In Section 6 we discuss our results, and finally certainly dominated by jets within several degrees to our line summarize our conclusions in Section 7. of sight, the radio–X-ray variability can result from changes in a variety of physical parameters: the synchrotron-emitting 2. VARIABILITY SIGNATURES OF DISKS AND JETS population of particles; the conditions under which they move; 2.1. Emission from Accretion Disks the orientation to line of sight; or the Doppler factor, δ ≡ [γ (1 − β cos θ)]−1, with γ = (1 − V 2/c2)−1/2 the Lorentz Fluctuations emerging directly from the surfaces of ADs factor, β = V/cthe velocity of the shock through the jet, and θ or from coronae above them can produce rest frame varia- the angle between the observer’s line of sight and the jet axis. 3 tions no faster than about tvar GMBH/c or 8 minutes for Major flares arise from new shocks passing through the jets = 8 MBH 10 M. The fastest real variations can provide a way (Marscher & Gear 1985), but smaller fluctuations could arise of measuring the lower limits on masses of the SMBHs. Most from turbulence behind those shocks (Marscher et al. 2008) X-ray and optical variability of the best studied cases, Seyfert or “mini-jets” (Giannios et al. 2009, 2010; Nalewajko et al. galaxies and radio-quiet quasars, is “red noise,” i.e., where the 2011). Alternatively, variations in the magnetic fields, changes PSD arising from the Fourier transform of the light curve is in the ambient medium, and changes in the injected particle ∝ α characterized by P (f ) f , with α<0, below some break distribution can also cause variations in emitted light. The light frequency, fb; above fb the PSD is usually dominated by mea- = travel time across the radiating region, modified by the Doppler surement errors with a white noise character (α 0) or perhaps factor, is the characteristic time scale. If these variable emissions by white noise with an astrophysical origin. The PSD can also arise from shocks in relativistic jets, as is expected to be the case yield the size of the largest emitting region if there is a flatter Δ −1 ∼− ∼− for FSRQs in high states, then the observed tvar R(δc) slope, often α 1 at low frequencies and 2 at intermediate (Gopal-Krishna et al. 2003). Here ΔR is the physical size of frequencies (e.g., Markowitz et al. 2003) before perhaps turning ∼ the emitting region in its rest frame (roughly a jet diameter) to 0 at the highest frequencies. Variations yielding red-noise and δ the Doppler factor. In this case, variability timescales can can be produced if energy is released over a wide range of time- constrain the jet’s δ (Jorstad et al. 2005). Constraints on the scales characteristic of different orbital periods and turbulent jet velocity and viewing angle can be particularly tight if other transport in ADs (e.g., Mangalam & Wiita 1993). data, such as apparent superluminal motions of radio knots, = − −1 2.2. Previous Indications of Quasi-periodic Variations where Vapp V sin θ(1 β cos θ) , can be detected through very long baseline interferometry. Our four FSRQ targets are Recently, detections of QPOs in five AGN have been made, all Very Long Baseline Array (VLBA) calibrators with compact mostly in X-rays. The best case is that of the Narrow Line Seyfert radio structures (e.g., [HB89] 1924+507=4C50.47, a member of 1 galaxy RE J1034+396 (Gierlinski´ et al. 2008), where a period the Caltech–Jodrell Bank flat-spectrum sample survey in Britzen of a little more than 1 hr was detected. The presence of optical et al. 2008).

ERRATUM: “AN X-RAYS SURVEY OF THE YOUNG STELLAR POPULATION OF THE LYNDS 1641 AND IOTA ORIONIS REGIONS” (2013, ApJ, 768, 99)

I. Pillitteri1,S.J.Wolk1, S. T. Megeath2, L. Allen3, J. Bally4, Marc Gagne´5,R.A.Gutermuth6, L. Hartman7, G. Micela8, P. Myers1, J. M. Oliveira9, S. Sciortino8, F. Walter1, L. Rebull10, and J. Stauffer10 1 SAO–Harvard Center for Astrophysics, 60 Garden St, Cambridge MA 02138, USA; [email protected] 2 Department of Physics & Astronomy, University of Toledo, OH, USA 3 National Optical Astronomy Observatory, USA 4 University of Colorado, Boulder, CO, USA 5 Department of Geology & Astronomy, West Chester University, West Chester, PA, USA 6 Department of Astronomy, University of Massachusetts, Amherst, MA 01003, USA 7 University of Michigan, Ann Arbor, MI, USA 8 INAF - Osservatorio Astronomico di Palermo, Italy 9 School of Physical & Geographical Sciences, Lennard-Jones Laboratories, Keele University, Staffordshire ST5 5BG, UK 10 CALTECH, Pasadena, CA, 91125, USA Received 2013 June 27; published 2013 July 29

Online-only material: color ﬁgure, machine-readable table

Due to an error at the publisher, the X-ray source identifiers in the published version of Table 4 are incorrect, which was brought to our attention by the HEASARC staff. We provide the revised table in electronic format with correct identifiers. The positional matches between the IR catalog and the X-ray catalog are corrected. The mistaken identifiers have no consequences for the published results. The same table in the online version is repeated twice. IOP Publishing sincerely regrets this error. In addition, the bottom left panel of Figure 3 has incorrect axis labels. We show here the panel plot with correct labels.

Table 4 List of X-Ray and IR Matches

SOXS ID R.A. Decl. JHK[3.6] [4.5] [5.8] [8.0] [24] 0/I/ﬂat II III r deg (J2000) deg (J2000) mag 1 83.4561 −6.3019 15.41 14.54 14.41 14.25 14.25 14.3 14.39 ... FFF1.4 2 83.4596 −6.3586 11.75 11.03 10.86 10.75 10.74 10.65 10.62 ... FFT1.3 3 83.4725 −6.2861 13.1 12.41 12.21 11.97 11.92 11.82 11.87 ... FFT1.5 4 83.4939 −6.2542 11.67 11.31 11.26 11.32 11.28 11.24 11.25 ... FFT1.4 5 83.5002 −6.5208 12.8 12.2 12.01 11.87 11.91 11.78 11.82 ... FFT0.9 6 83.5087 −6.4517 ...... 16.7 15.73 15.63 13.95 ... FFF1.3 6 83.5084 −6.4528 ...... 16.85 ...... FFF3.3 7 83.5159 −6.2679 12.04 11.36 11.18 10.98 10.92 10.86 10.84 ... FFT1.4 8 83.5167 −6.2249 12.7 12.11 11.91 11.7 11.68 11.62 11.61 ... FFT0.6 9 83.5259 −6.5125 11.59 10.84 10.59 10.42 10.39 10.34 10.3 5.68 F T T 2.1 10 83.5278 −6.4033 11.68 10.83 10.51 10.43 10.4 10.3 10.26 ... FFT1.7

Notes. Columns are: SOXS source number, IR coordinates, 2MASS J, H, K magnitudes, IRAC [3.6], [4.5], [5.8], [8.0], MIPS [24] band magnitudes, a ﬂag indicating Class 0/I/ﬂat-spectrum object, Class II object, or Class III candidate status, and the match radius between X-ray and IR positions. X-ray sources with more than a potential IR counterpart have multiple entries. (This table is available in its entirety in a machine-readable form in the online journal. A portion is shown here for guidance regarding its form and content.)

1 The Astrophysical Journal, 773:80 (2pp), 2013 August 10 Erratum: 2013,ApJ,768, 99

Figure 3. IRAC [3.6] − [5.8] vs. [4.5] − [8.0]. See the caption of Figure 3 in the published version of the paper. (A color version of this ﬁgure is available in the online journal.)

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Friction forces on phase transition fronts

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Please note that terms and conditions apply. JCAP07(2013)045 hysics P le ic opic interaction of the t ogy connection, primor- and, more recently, ultra- ar 10.1088/1475-7516/2013/07/045 en study the velocity of phase ultra-relativistic velocities and ion, which interpolates between compute these forces, and only the thermodynamic parameters, ons for the fronts to actually run doi: ers allow the existence of runaway cles manifests itself macroscopically strop , the friction depends on two friction A osmology and and osmology C 1 1303.4233 cosmological phase transitions, particle physics - cosmol

In cosmological ﬁrst-order phase transitions, the microsc [email protected] rnal of rnal

ou An IOP and SISSA journal An IOP and Member of CONICET, Argentina. 1 2013 IOP Publishing Ltd and Sissa Medialab srl IFIMAR (CONICET–UNMdP), Departamento deFacultad F´ısica, de Ciencias Exactas FunesDe´an y 3350, Naturales, (7600) UNMdP, Mar del Plata, Argentina E-mail: c dial gravitational waves (theory) ArXiv ePrint:

Received April 2, 2013 Accepted July 8, 2013 Published July 29, 2013 Abstract. phase transition fronts with non-equilibriumas plasma parti friction forces.two In limits general, have itrelativistic been is walls studied, a which nontrivial namely, runshow problem away. that that stationary to In of solutions this still verywalls. paper exist Hence, slow when we we the walls consider discuss paramet away. the We necessary also and propose sufficient a conditi the phenomenological model non-relativistic for and the frict ultra-relativisticcoefficients values. which can Thus betransition calculated for fronts specific as models.and a We the function th amount of of the supercooling. friction parameters, Keywords: Ariel Mégevand Friction forces on phasefronts transition J The Astrophysical Journal, 773:95 (5pp), 2013 August 20 doi:10.1088/0004-637X/773/2/95 C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

SOLUBILITY OF IRON IN METALLIC HYDROGEN AND STABILITY OF DENSE CORES IN GIANT PLANETS

Sean M. Wahl1,HughF.Wilson1,2, and Burkhard Militzer1,3 1 Department of Earth and Planetary Science, University of California, Berkeley, CA 94720, USA 2 Virtual Nanoscience Laboratory, CSIRO Materials Science and Engineering, Parkville, Victoria 3052, Australia 3 Department of Astronomy, University of California, Berkeley, CA 94720, USA Received 2013 March 26; accepted 2013 June 17; published 2013 July 29

ABSTRACT The formation of the giant planets in our solar system, and likely a majority of giant exoplanets, is most commonly explained by the accretion of nebular hydrogen and helium onto a large core of terrestrial-like composition. The fate of this core has important consequences for the evolution of the interior structure of the planet. It has recently been shown that H2O, MgO, and SiO2 dissolve in liquid metallic hydrogen at high temperature and pressure. In this study, we perform ab initio calculations to study the solubility of an innermost metallic core. We find dissolution of iron to be strongly favored above 2000 K over the entire pressure range (0.4–4 TPa) considered. We compare with and summarize the results for solubilities on other probable core constituents. The calculations imply that giant planet cores are in thermodynamic disequilibrium with surrounding layers, promoting erosion and redistribution of heavy elements. Differences in solubility behavior between iron and rock may influence evolution of interiors, particularly for Saturn-mass planets. Understanding the distribution of iron and other heavy elements in gas giants may be relevant in understanding mass–radius relationships, as well as deviations in transport properties from pure hydrogen–helium mixtures. Key words: planets and satellites: dynamical evolution and stability – planets and satellites: individual (Jupiter, Saturn) – planets and satellites: interiors Online-only material: color figures

1. INTRODUCTION As a result of continuing discoveries by Kepler (Borucki et al. 2010) and other exoplanet surveys, the number of confirmed Despite recent advances in computational methods improving planets has climbed to over 800, the majority of which are gi- understanding of the hydrogen–helium dominated outer layers ants. This presents a growing sampling of planetary mass–radii (McMahon et al. 2012; French et al. 2012; Militzer 2013; Wilson relationships that will be fundamental to understanding the evo- & Militzer 2010), knowledge of the deep interior structure lution of giant planet interiors. The range of mass–radius re- of giant planets is limited. Determining the size of a dense lationships observed for exoplanets exhibit variation beyond central core in a giant planet is dependent upon the model and those in the solar system. In some cases, such as Corot-20b equation of state used. Current observational evidence yields (Deleuil et al. 2011), the relationships may even defy explana- recent estimates for present day core mass of ∼0–10 (Guillot tion by simple structural models. Redistribution of dense core 2005), ∼14–18 (Militzer et al. 2008) Earth masses for Jupiter, material lowers the heavy element content required to explain and ∼9–22 (Guillot 2005) Earth masses for Saturn. The Juno anomalously high observed densities. spacecraft, en route to Jupiter, will improve this constraint The favored model for gas giant formation (Mizuno et al. with more precise measurements of the giant’s gravitational 1978; Bodenheimer & Pollack 1986; Pollack et al. 1996) field (Helled 2011). Meanwhile, the density profiles of Neptune relies on the early formation of a large planetary embryo and Uranus allow non-unique solutions for the compositional of critical mass to cause runaway accretion of hydrogen and structure for much of the interior (Guillot 1999, 2005). helium gas. A competing theory involves collapse of a region It has long been suggested (Stevenson 1982a, 1982b) that of the disk under self-gravity, e.g., (Boss 1997), but may a portion of this dense material might be redistributed in have difficulty explaining significant enrichment of refractory solution with hydrogen. As a result, erosion of a dense core elements (Hubbard et al. 2002; Guillot 2005). The immediate would cause it to shrink over the lifetime of the planet. result of a core-accretion hypothesis is a planet with the Possible consequences of this process are only beginning to be ice–rock–metal embryo residing at the center as a dense core, enumerated in evolutionary models (Chabrier & Baraffe 2007; surrounded by an extensive layer of metallic hydrogen and Leconte & Chabrier 2012; Mirouh 2012). The establishment helium. The role of core erosion to the subsequent evolution of a gradient in concentration of a heavy dissolved component is a major source of uncertainty, but, in principle, can explain may change the nature of convection in a portion of the planets shrinking of cores to masses smaller than those necessary to interior. This “double-diffusive” is hypothesized to reduce the form the planet under the core-accretion hypothesis. efficiency of heat transfer, thereby altering the thermal evolution Core erosion in giant planets can be addressed by determin- of the planet. Comprehensive understanding of the process has ing the solubility of analogous phases. Previous studies have been limited by the lack of knowledge of the solubility of various considered an icy layer of fluid and superionic H2O (Wilson & phases in metallic hydrogen, as well as poor understanding of Militzer 2012a; Wilson et al. 2013), and a rocky layer consist- the scaling of convective efficiency in the presence of competing ing of MgO (Wilson & Militzer 2012b) and SiO2 (Gonzalez gradients of composition and temperature. In this study, we et al. 2013), which have been shown to separate at rele- address the first issue for iron metal. vant conditions (Umemoto 2006). Assuming the same gross

1 The Astrophysical Journal, 773:95 (5pp), 2013 August 20 Wahl, Wilson, & Militzer distribution, elements as terrestrial bodies, the innermost core where the averaging is over configurations generated dur- would be composed of a dense, metallic alloy composed pri- ing simulations of the system governed by the hybrid marily of iron. potential. Ab initio random structure searches (Pickard & Needs 2009) To increase the efficiency, we calculated the Helmholtz free demonstrate that iron remains in a hexagonal close packed (hcp) energy in two steps each involving an integral of the form of structure remains stable up to pressures approaching Jupiter’s Equation (2). We first find ΔFDFT→cl, between the systems gov- center, ∼2.3 TPa, at which point it undergoes a phase transition erned by DFT and classical pair potentials, which are fit via to a face centered cubic (fcc) structure. Stixrude (2012) demon- a force-matching method (Izvekov et al. 2004). In the second strated a gradual decrease in this transition pressure with tem- step, we find ΔFcl→an, between the pair potentials and a refer- perature. Simulations of liquid hydrogen (Militzer et al. 2008; ence system with an analytic solution Fan. The Helmholtz en- Militzer 2013; McMahon et al. 2012) undergo a gradual transi- ergy for the DFT systems, FDFT = Fan + ΔFDFT→cl + ΔFcl→an, tion from molecular to metallic, which is complete by ∼0.4 TPa may then be compared directly. The first step requires DFT-MD at low temperatures. Stable mixtures of Fe and H have been sug- simulations for a small number of λ values, while the second gested at lower P–T conditions, applicable to terrestrial cores applies a much faster classical Monte Carlo approach at numer- (Bazhanova 2012). ous values of λ to ensure a smooth integration. Finding a suitable reference system is essential to the method, Δ 2. THEORY AND METHODOLOGY as it allows Fcl→DFT to be found with a small number of number of integration steps, and prevents solids from melting or Compression of materials with modern experimental transforming to a new structure. For liquid systems, we integrate techniques can reach megabar pressures, however, the to a non-interacting ideal gas, using classical two-body pair pressure–temperature conditions near Jupiter’s core (>4TPa) potentials as the intermediate step. For solid iron, we integrate to remain inaccessible. Temperatures in shockwave experiments an Einstein solid, a system of non-interacting, three-dimensional climb rapidly at high pressures in comparison with planetary harmonic oscillators. For the intermediate classical solid system adiabats, while diamond anvil cell experiments are best suited we use a 50–50 “mixture” of two-body and one-body harmonic for low temperatures. As a result, simulations based on ab initio oscillator potentials. Spring constants for the Einstein terms are theories are best suited for directly probing the conditions of found from the mean-squared displacement of atoms from their gas giant interiors. ideal lattice sites during a DFT-MD simulation. We performed density functional theory molecular dynamics All simulations presented here were performed using the (DFT-MD) simulations to determine the energetics of a dis- Vienna ab initio simulation package (VASP; Kresse 1996). solution reaction, in which solid iron dissolves in pure liquid VASP uses the DFT formalism utilizing pseudopotentials of the hydrogen. We calculate a Gibbs free energy of solvation: projector augmented wave type (Blochl 1994) and the exchange- correlation functional of Perdew et al. (1996). The iron pseu- 1 dopotential treats an [Mg]3pd64s2 electron configuration as ΔG (Fe : 256H) = G (H Fe) − G (H ) + G (Fe ) , sol 256 256 32 32 valence states, and a 2 × 2 × 2gridofk-points is used for (1) all simulations. Simulations on hydrogen and the solution were performed with a 900 eV cutoff energy for the plane wave expansion, while a 300 eV cutoff was used for iron. A time where G (H256) and G (Fe32) are the Gibbs free energies of a pure hydrogen liquid and solid or liquid iron. G (H Fe) step of 0.2 fs was used for all liquid simulations, a 0.5–1.0 fs 256 time step was used for high- and low-temperature iron simu- is the Gibbs free energy of 1:256 liquid solution of iron in Δ hydrogen. We assume that analysis of a single low-concentration lations, respectively. The Gsol results were confirmed to be solution is sufficient to determine the onset of core erosion, well converged with respect to the energy cutoff and time step. since the reservoir of metallic hydrogen would be much larger Prohibitively long simulation times required that convergence with respect to finer k-point meshes be verified over a subset than the core. This does not rule out non-ideal effects of higher × × concentrations that might exist in a narrow, poorly convecting of configurations generated by a simulation with a 2 2 layer at the top of a core. 2grid. 2.2. Material Phases 2.1. Computation of Gibbs Free Energies Iron simulations assume an hcp or fcc structure within their Free energy calculations require the determination of a con- respective stability regimes (Pickard & Needs 2009; Stixrude tribution from entropy, which is not determined from the stan- 2012). We confirmed Fe to be solid up to 20,000 K at 4 TPa, dard DFT-MD formalism. To achieve this, we adopt a two-step and to be a liquid at temperatures as low as 15,000 K at thermodynamic integration method, as used in previous studies 1 TPa. We also confirmed that the Gibbs free energy favors hcp (Morales et al. 2009; Wilson & Militzer 2010, 2012a, 2012b). stability over fcc at 1 TPa, though the difference is negligible The method requires integration of the change in Helmholtz for our subsequent analysis of dissolution. We found 32 atom free energy over an unphysical, yet thermodynamically per- supercells to be sufficient for Fe simulation. Finite size effects missible, transformation between two systems governed by required that we use large 256 atom supercells for hydrogen, to potentials Ua (ri) and Ub (ri). We define a hybrid potential which one Fe atom was added for the solution. Cubic supercells = − Uλ (1 λ) Ua + λUb, where λ is the fraction of the po- are used for fcc and liquid runs. In order to maintain the tential Ub (ri). The difference is Helmholtz free energy is then same number of atoms for the hcp an orthogonal supercell given by defined the combination of hexagonal unit cell vectors a, a + b, and c. 1 Cell volumes at each temperature were determined by fit- ΔFb→a ≡ Fb − Fa = dλUb (ri) − Ua (ri)λ, (2) 0 ting a pressure–volume polytrope equation of state to short

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Non-maximal , large and tri-bimaximal via quark-lepton complementarity at next-to- leading order

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Non-maximal θ23,largeθ13 and tri-bimaximal θ12 via quark-lepton complementarity at next-to-leading order

J. Harada

Niihama National College of Technology - Niihama 792-8580, Japan

received 11 May 2013; accepted in ﬁnal form 5 July 2013 published online 29 July 2013

PACS 14.60.Pq – Neutrino mass and mixing

Abstract – We present analytical formulae for the neutrino mixing angles at the next-to-leading order in the quark-lepton complementarity, and show that higher-order corrections are important to explain the observed pattern of neutrino mixing. In particular, the next-to-leading–order corrections 1) lead to a deviation of θ23 from maximal mixing, 2) reduce the predicted value 2 2 of sin 2θ13 by 9.8%, 3) provide the same value of sin θ12 as that of the tri-bimaximal mixing.

√ ◦ ◦ Introduction and motivation. – The main recent These correspond to θ13 θC/ 2 9 , θ23 45 and ◦ developments on neutrino mixing [1–5] are related to the θ12 36 , which are consistent with experimental results relatively large value of θ13, which was measured by recent at leading-order approximation. In particular, it is inter- ◦ experiments [6–10], and indications of significant deviation esting that the predicted value of 1-3 mixing, θ13 9 , of θ23 from maximal mixing [11,12]. In particular, non- has been confirmed by recent experiments [6–10]. maximal θ23 is strongly indicated by recent data, but the However, the next-to-leading–order corrections, the 4 sign of θ23 − π/4 is not yet determined. Results of global O(λ ) terms in eqs. (2)–(4), should be calculated, because: analysis are given in refs. [13–15]. –theO(λ4) corrections may not be enough small with To explain the observed pattern of neutrino mixing, the respect to the recent experimental errors; particu- quark-lepton complementarity (QLC) has been widely in- larly, the Daya Bay Collaboration reported the precise vestigated in the literature [16–42]. In particular, much value at 1σ [9], attention has been paid to the class of models, UPMNS = † V V 2 CKM M , which can be obtained in grand unified theories sin 2θ13 =0.092 ± 0.016(stat) ± 0.005(syst), (5) (GUTs). The correlation matrix VM is simply defined by the product of the CKM and PMNS mixing matrices. In where the magnitude of systematic errors is the same 4 general VM is not determined by theories, because there order of that of λ ; is no relation between the Dirac and the Majorana mass ◦ – the same value of 1-3 mixing, θ13 9 ,canbe operators [26]. In this class of models, the observed two obtained in various models with different schemes: large mixing angles θ12 and θ23 indicate that VM has two flavor symmetries, texture, ansätz etc. [12]; to dis- large mixing angles because all the mixing angles in VCKM tinguish models, we need a more precise prediction; are small. As the simplest possibility, for example, we can V V take M being the bimaximal mixing matrix bm, –a deviationof θ23 from maximal mixing is the O(λ4) correction in eq. (3); to determine the mag- U V † V . PMNS = CKM bm (1) nitude of deviation, the O(λ4) corrections should be In this paper we only consider this minimal model. In calculated; 4 ref. [26], expanding VCKM to O(λ )(whereλ ≡ sin θC ≈ – in the earlier works, eq. (1) was analyzed numerically; 0.2253), it is found that analytical formulae of higher-order terms are usually 2 2 4 4 neglected; therefore, it is very unclear that which pa- sin 2θ13 =2λ + O(λ )=0.102 + O(λ ), (2) rameter is relevant at each order. 2 4 sin 2θ23 =1+O(λ ), (3) 2 2 4 4 Motivated by these points, in this paper we perform θ − λ O λ . O λ . 4 sin 2 12 =1 2 + ( )=0898 + ( ) (4) analytical calculations of the O(λ ) corrections. We show

KECK OBSERVATIONS OF THE GALACTIC CENTER SOURCE G2: GAS CLOUD OR STAR?

K. Phifer1,6,T.Do2,7, L. Meyer1,A.M.Ghez1,8, G. Witzel1, S. Yelda1, A. Boehle1,9, J. R. Lu3,10, M. R. Morris1, E. E. Becklin1,4, and K. Matthews5 1 Department of Physics and Astronomy, University of California, Los Angeles, CA 90095, USA; [email protected] 2 Dunlap Institute for Astronomy and Astrophysics, University of Toronto, Toronto, Ontario, Canada M5S 3H4 3 Institute for Astronomy, University of Hawaii, Hilo, HI 96720, USA 4 NASA-Ames Research Center, Mountain View, CA 94035, USA 5 Division of Physics, Mathematics, and Astronomy, California Institute of Technology, Pasadena, CA 91125, USA Received 2013 April 17; accepted 2013 July 7; published 2013 July 29

ABSTRACT We present new observations and analysis of G2—the intriguing red emission-line object which is quickly approaching the Galaxy’s central black hole. The observations were obtained with the laser guide star adaptive optics systems on the W. M. Keck I and II telescopes (2006–2012) and include spectroscopy (R ∼ 3600) centered on the hydrogen Brγ line as well as K (2.1 μm) and L (3.8 μm) imaging. Analysis of these observations shows the Brγ line emission has a positional offset from the L continuum. This offset is likely due to background source confusion at L. We therefore present the ﬁrst orbital solution derived from Brγ line astrometry, which, when coupled with radial velocity measurements, results in a later time of closest approach (2014.21 ± 0.14), closer periastron (130 AU, 1600 Rs), and higher eccentricity (0.9814±0.0060) compared to a solution using L astrometry. It is shown that G2 has no K counterpart down to K ∼ 20 mag. G2’s L continuum and the Brγ line emission appears unresolved in almost all epochs, which implies that the bulk of the emission resides in a compact region. The observations altogether suggest that while G2 has a gaseous component that is tidally interacting with the central black hole, there is likely a central star providing the self-gravity necessary to sustain the compact nature of this object. Key words: accretion, accretion disks – black hole physics – Galaxy: center – Galaxy: kinematics and dynamics – infrared: general

1. INTRODUCTION Murray-Clay & Loeb 2012; Scoville & Burkert 2013). In these scenarios, the observed heated gas is posited to be circumstellar, Recently, Gillessen et al. (2012, 2013) reported the discovery either intrinsic or a consequence of interaction of the star and of G2, an extremely red object with spatially resolved Brγ surrounding ambient gas. In these stellar scenarios, G2 not emission. The object was interpreted as a ∼3 Earth-mass gas only has existed for timescales much longer than the observed cloud based upon an inferred low dust temperature, observed time baseline, but the continuum source (L) will also survive elongation of the Brγ emission along the object’s direction of periapse passage. motion, and a claimed tail along the same orbital trajectory as Regardless of its nature, G2’s properties and possible origin G2. This interpretation is particularly interesting because G2 depend critically on its orbital parameters. These parameters is on a highly eccentric orbit with a closest approach to our have been estimated from observations with a short time Galaxy’s central black hole within the next year, potentially baseline compared to the orbital period (∼10 versus ∼200 yr; allowing us to observe an unprecedented accretion event onto Gillessen et al. 2013) and in a very crowded region, making a supermassive black hole and offer insight into the region the orbital solution susceptible to biases (Hartkopf et al. 2001). surrounding the black hole (e.g., Morris 2012;Moscibrodzka´ We therefore present new measurements and analysis of G2 that et al. 2012; Anninos et al. 2012; Saitoh et al. 2012; Bartos et al. minimize the effects of source confusion on estimates of G2’s 2013; Yusef-Zadeh & Wardle 2013). orbital parameters and examine the temporal evolution of G2’s The interpretation of G2 as a gas cloud, however, is not properties. definitive. One challenge for the pure gas cloud scenario is that, given the strong tidal fields in this region and G2’s low self- 2. OBSERVATIONS gravity, G2 must have formed quite recently (∼1995, just prior Two types of new data were collected for this study using the to the initial observations; Burkert et al. 2012; Schartmann et al. laser guide star adaptive optics (LGS AO) systems at the W.M. 2012). Since such a gas cloud would be tidally disrupted during Keck Observatory (Wizinowich et al. 2006;vanDametal.2006; its periapse passage in the upcoming year, the gas cloud model Chin et al. 2012). Spectra were obtained using the OSIRIS inte- implies that G2 will be observed over almost the entire extent of gral field spectrograph (Larkin et al. 2006) through the narrow- its existence. Therefore, several alternative scenarios invoking band Kn3 filter, which is centered on the Brγ hydrogen line an underlying star have been proposed (Miralda-Escude´ 2012; (2.1661 μm), at a spectral resolution of R ∼ 3600. Imaging data were obtained in the K -band filter (λ0 = 2.124 μm) and 6 NSF Graduate Student Fellow, Grant DGE-0707424. the L -band filter (λ0 = 3.776 μm) using the Keck II near- 7 Dunlap Fellow. 8 Leichtman-Levine Chair in Astrophysics. infrared camera (NIRC2; PI: K. Matthews). These data were 9 Preston Fellow. obtained and reduced in a similar manner as in our previous 10 NSF Astronomy and Astrophysics Postdoctoral Fellow, publications (Ghez et al. 2005b; Hornstein et al. 2007; Ghez Grant AST-1102791. et al. 2008;Doetal.2009;Luetal.2009; Yelda et al. 2010;

1 The Astrophysical Journal Letters, 773:L13 (6pp), 2013 August 10 Phifer et al.

Table 1 Summary of Observations and Measurements of G2

b UT Date Fractional AO Type/ Pixel Scale Nframes Nframes FWHM Orig vlsr Brγ FWHM G2 ΔR.A. G2 ΔDecl. Date Telescope (mas) Observed Used (mas) Pub.a (km s−1)(kms−1) (mas) (mas) OSIRIS, Kn3 2006 Jun 18, 30; Jul 1 2006.495 Keck II LGS 35 28 27 74 3 1125 ± 6 137 ± 16 222.97 ± 5.05 −105.40 ± 2.40 2008 May 16; Jul 25 2008.487 Keck II LGS 35 22 21 78 4,0 ee181.82 ± 6.90 −64.97 ± 2.51 2009 May 5, 6 2009.344 Keck II LGS 35 24 19 79 0 1352 ± 20 163 ± 28 176.26 ± 2.74 −65.76 ± 2.04 2010 May 5, 8 2010.349 Keck II LGS 35 17 16 82 0 1479 ± 18 256 ± 33 166.29 ± 9.11 −43.95 ± 11.05 2012 Jun 9, 11; Aug 11, 12 2012.613 Keck I LGS 20c 27 21d 68 0 2071 ± 86 726 ± 111 103.16 ± 4.07 −8.79 ± 11.36 NIRC2, L 2003 Jun 10 2003.440 Keck II NGS 10 12 12 86 0 260.1 ± 4.5 −154.6 ± 5.1 2004 Jul 26 2004.567 Keck II LGS 10 11 11 80 1 262.6 ± 4.8 −140.4 ± 4.2 2005 Jul 30 2005.580 Keck II LGS 10 62 56 87 2 251.6 ± 3.4 −129.0 ± 2.6 2006 May 21 2006.385 Keck II LGS 10 19 19 81 0 226.2 ± 1.9 −99.4 ± 0.8 2009 Jul 22 2009.561 Keck II LGS 10 4 4 85 0 175.4 ± 4.9 −70.7 ± 2.1 2012 Jul 20–23 2012.562 Keck II LGS 10 1316 1314 93 0 105.9 ± 0.8 −21.1 ± 0.6

Notes. a References: (0) this Work; (1) Ghez et al. (2005a); (2) Hornstein et al. (2007); (3) Ghez et al. (2008); (4) Do et al. (2009). b Offset from SgrA*-radio. c A smaller square dither pattern of 0. 5 × 0. 5 was used. d In 2012, a more stringent quality cut (FWHM < 68 mas) was used because G2 is closer to stars than in other epochs and because the improved LGS performance on Keck I allowed this more stringent cut. e In 2008, low signal-to-noise ratio and changes in the local standard of rest velocity between the two observation dates prevent reliable line measurements.

Meyer et al. 2012) and the speciﬁc spectroscopic and L observations utilized in this Letter are described in 2012.613 Table 1.

3. ANALYSIS 2010.349 3.1. OSIRIS IFU Measurements Because G2 has no detectable K continuum and is fainter rest than any object we have previously extracted, some analysis γ steps differ from our earlier analyses. For all epochs, we created 2009.344 Br− a combined data cube before extracting G2’s spectrum rather Arbitrary Flux Units than extracting spectra from individual cubes. The OH sky lines in the data are subtracted using sky frames scaled to the strength of families of OH lines in the observed frames to account for 2006.495 temporal variations. In 2012, the Brγ emission line from G2 is coincident with a prominent OH sky line at 2.180 μm, so 2.14 2.16 2.18 2.20 2.22 μ to minimize the systematic effects associated with OH line Wavelength ( m) subtraction, we scale the sky only to this line. Figure 1. The evolving spectra of G2. The highly redshifted Brγ emission line An iterative process was required to estimate the position has a FWHM that increases with time (see Table 1). The Brγ line peak has been scaled such that the flux scales are constant. Additionally, the spectra have and spectral properties of G2 in the OSIRIS data cubes (see been smoothed using a boxcar average with a width of three OSIRIS channels Figures 1 and 2) since G2’s position is needed to place the (7.5 × 10−4 μm). aperture for spectral extraction, and G2’s spectral properties are needed to determine which OSIRIS channels should be used to measure its position. G2’s position was first obtained by and spectrum for G2. Further iterations produced no significant visual inspection of the data cube. Then, an initial spectrum changes. Measuring G2’s position on a continuum-subtracted was extracted at this position using an aperture with a radius frame (see Figure 2(b)) removes the effect of source confusion, of 35 mas. Emission from local ambient gas was subtracted thereby avoiding astrometric biases. using a region free of stellar halos within ∼0.5. A Gaussian G2’s spectral properties were obtained using the final Gaus- was fit to the resulting emission line. In order to refine the sian fits to the Brγ emission line. The radial velocities (RVs) position, the three-dimensional data cube was median collapsed for G2 were calculated from the offset of the line from the over the wavelength range corresponding to twice the standard rest wavelength (λBrγ = 2.1661 μm) and corrected to the local deviation of the Gaussian fit to the emission line, centered on standard of rest. The reported FWHM measurements were cor- the line peak. Continuum emission was subtracted by averaging rected for instrumental broadening (FWHM = 85 km s−1). Line the median of the 25 nearest channels on either side of the flux measurements were made by comparing the integral of the line. G2’s position was further refined with a two-dimensional Gaussian fit to the integral of the flux density of the non-variable Gaussian fit to the continuum-subtracted image. The iterative star S0-2 (Fobserved,K = 1.46±0.02 mJy; Ghez et al. 2008) over extraction process was repeated again to obtain a final position the wavelength range of 2.17–2.18 μm. Dereddened fluxes were

GMRT DISCOVERY OF PSR J1544+4937: AN ECLIPSING BLACK-WIDOW PULSAR IDENTIFIED WITH A FERMI-LAT SOURCE

B. Bhattacharyya1,J.Roy1,P.S.Ray2, Y. Gupta1, D. Bhattacharya3,R.W.Romani4,S.M.Ransom5, E. C. Ferrara6, M. T. Wolff2,F.Camilo7,8, I. Cognard9,10, A. K. Harding6, P. R. den Hartog4, S. Johnston11, M. Keith11, M. Kerr4, P. F. Michelson4, P. M. Saz Parkinson12,D.L.Wood13,14, and K. S. Wood2 1 National Centre for Radio Astrophysics, Tata Institute of Fundamental Research, Pune 411 007, India 2 Space Science Division, Naval Research Laboratory, Washington, DC 20375-5352, USA 3 Inter-University Centre for Astronomy and Astrophysics, Pune 411 007, India 4 W. W. Hansen Experimental Physics Laboratory, Kavli Institute for Particle Astrophysics and Cosmology, Department of Physics and SLAC National Accelerator Laboratory, Stanford University, Stanford, CA 94305, USA 5 National Radio Astronomy Observatory (NRAO), Charlottesville, VA 22903, USA 6 NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA 7 Columbia Astrophysics Laboratory, Columbia University, New York, NY 10027, USA 8 Arecibo Observatory, Arecibo, Puerto Rico, PR 00612, USA 9 Laboratoire de Physique et Chimie de l’Environnement, LPCE UMR 6115 CNRS, F-45071 Orleans´ Cedex 02, France 10 Station de radioastronomie de Nan¸cay, Observatoire de Paris, CNRS/INSU, F-18330 Nan¸cay, France 11 CSIRO Astronomy and Space Science, Australia Telescope National Facility, Epping, NSW 1710, Australia 12 Santa Cruz Institute for Particle Physics, Department of Physics and Department of Astronomy and Astrophysics, University of California at Santa Cruz, Santa Cruz, CA 95064, USA 13 Praxis Inc., Alexandria, VA 22303, USA 14 Naval Research Laboratory, Washington, DC 20375, USA Received 2013 April 26; accepted 2013 June 24; published 2013 July 29

ABSTRACT Using the Giant Metrewave Radio Telescope, we performed deep observations to search for radio pulsations in the directions of unidentified Fermi-Large Area Telescope γ -ray sources. We report the discovery of an eclipsing black- widow millisecond pulsar, PSR J1544+4937, identified with the uncataloged γ -ray source Fermi J1544.2+4941. This 2.16 ms pulsar is in a 2.9 hr compact circular orbit with a very low mass companion (Mc > 0.017M). At 322 MHz this pulsar is found to be eclipsing for 13% of its orbit, whereas at 607 MHz the pulsar is detected throughout the low-frequency eclipse phase. Variations in the eclipse ingress phase are observed, indicating a clumpy and variable eclipsing medium. Moreover, additional short-duration absorption events are observed around the eclipse boundaries. Using the radio timing ephemeris we were able to detect γ -ray pulsations from this pulsar, confirming it as the source powering the γ -ray emission. Key words: binaries: eclipsing – pulsars: general – pulsars: individual (PSR J1544+4937) Online-only material: color figures

1. INTRODUCTION mass companion (Mc 0.1M) by the pulsar wind. There were two such eclipsing BW systems in the Galactic field known The Large Area Telescope (LAT; Atwood et al. 2009)on before the launch of Fermi—PSR B1957+20 (Fruchter et al. board the Fermi Gamma-ray Space Telescope has discovered 1988) and PSR J2051−0827 (Stappers et al. 1996). The BW a large number of γ -ray point sources, of which many are pulsars are found to have higher values of spin-down energy-loss unidentified or even unassociated with any known potential rate (E˙ ∼ 1034 erg s−1) compared to other MSPs, making these counterpart (Ackermann et al. 2012). The LATcan localize most systems good candidates for pulsed γ -ray emission (Roberts of these sources well enough that they can be covered in a single 2011). Among 43 new MSPs found in Fermi-directed searches, pointing with the large primary beams of radio telescopes at low there are at least 10 BWs (Ray et al. 2012). This Letter frequencies, allowing them to be searched efficiently. Targeted describes the discovery and follow-up study of an eclipsing BW radio searches of unassociated LAT point sources by the Fermi MSP, J1544+4937, with the Giant Metrewave Radio Telescope Pulsar Search Consortium (PSC) have resulted in the discovery (GMRT). of 43 radio millisecond pulsars (MSPs; Ray et al. 2012). MSPs are thought to evolve from normal pulsars in binary systems 2. OBSERVATIONS AND SEARCH ANALYSIS via transfer of angular momentum from companions. Thus, the majority of MSPs are naturally expected to be in binaries As a part of the PSC search effort, we observed mid- and high- (∼83% being the binary fraction for MSPs in the Galactic Galactic-latitude unassociated Fermi point sources with the field15). Binary systems where the pulsar wind evaporates the GMRT at 607 MHz. The GMRT Software Back-end (Roy et al. companion are one way to form isolated MSPs. Such systems 2010) produces simultaneous incoherent and coherent filter- bank outputs of 512 × 0.0651 MHz sampled every 61.44 μs. where the interaction is ongoing are called black-widow (BW) pulsars. Many exhibit long eclipses (∼10% of the orbital period, The wider incoherent beam of the GMRT (40 at 607 MHz) can apparently larger than the companion’s Roche lobe) that are easily cover error circles associated with the Fermi sources. In believed to be caused by the material blown from the very low addition, a coherent beam that is three times more sensitive and narrower (1.5 at 607 MHz using the central core of the GMRT) can be useful if the pulsar happens to be near the pointing 15 http://astro.phys.wvu.edu/GalacticMSPs/GalacticMSPs.txt center.

1 The Astrophysical Journal Letters, 773:L12 (6pp), 2013 August 10 Bhattacharyya et al.

Table 1 Parameters of J1544 + 4937 Parameter Valuea Interferometric positionb Right ascension (J2000) ... 15h44m04s.166 ± 0s.3 Declination (J2000) ... +49◦3757. 45 ± 4. 7 Offset from survey beam center ... 4.3 Parameters from radio timing Right ascension (J2000) ... 15h44m04s.48722 (2) Declination (J2000) ... +49◦3755. 2545 (2) Position epoch (MJD) ... 51544.0 Pulsar period, P (ms) ... 2.15928839043289 (5) Pulsar frequency, f (Hz) ... 463.11553585462 (1) Frequency derivative, f˙ (Hz s−1) ... −6.29 (1) × 10−16 Period epoch (MJD) ... 56007.0 Dispersion measure, DM (cm−3 pc) ... 23.2258 (11) Binary model ... ELL1 Figure 1. Post-fit timing residuals of J1544+4937 considering non-eclipsing binary phases. Orbital period, Pb (days) ... 0.1207729895 (1) Projected semi-major axis, x (lt-s) ... 0.0328680 (4) (A color version of this figure is available in the online journal.) Epoch of ascending node passage, TASC (MJD) ... 56124.7701121 (2) Span of timing data (MJD) ... 652 width, 10% duty-cycle, incoherent array gain of 2.3 K Jy−1,for Number of TOAs... 280 30 minutes of observing, we estimate the search sensitivity as Post-fit residual rms (μs) ... 6.9 (92 K + Tsky)/(335 K) mJy for a 5σ detection at 607 MHz. Con- Reduced chi-square ... 2.7 ◦ sidering |b| > 5 , where Tsky ∼ 10–45 K, our search sensitivity Derived parameters is 0.3–0.4 mJy. Mass function, f (M) ... 0.0000026132 In a 30 minute pointing on 2011 February 1, toward Fermi Min companion mass, mc (M) ... 0.017 J1544.2+4941 we discovered a binary MSP of period 2.16 ms DM distancec (kpc) ... 1.2 with significant acceleration of 2.25 m s−2 at a DM of Flux density at 322 MHz (mJy) ... 5.4 23.2pccm−3. Flux density at 607 MHz (mJy) ... 1.2 Spectral index ... −2.3 8 3. FOLLOW-UP TIMING Surface magnetic field, Bs (10 G) ... 0.805 (1) Spin-down luminosity, E˙ (1034 erg s−1) ... 1.150 (8) We localized J1544 + 4937 with an accuracy of 5 (positions Characteristic age, τ (Gyr) ... 11.65 (3) listed in Table 1) using continuum imaging for the full GMRT γ -ray parametersd array followed by multi-pixel beamforming (Roy et al. 2012), Photon flux (>0.1GeV,cm−2 s−1) ... 1.6(8) × 10−9 which allowed us to have sensitive follow-up studies using the Energy flux (>0.1 GeV, erg cm−2 s−1) ... 2.1(6) × 10−12 coherent array. We estimate a flux of 5.4 mJy at 322 MHz, −2 −1 32 − Luminosity, Lγ /f Ω (>0.1 GeV, erg cm s ) ... 3.6 × 10 and a spectral index of 2.3. We started the regular timing Efficiency, ηγ /f Ω (>0.1GeV)... 0.03 campaign for J1544 + 4937 in 2011 April at 322 MHz with the same coherent filter bank. With the derived position from Notes. the multi-pixel search and an a priori binary model predicted a Errors in the last digit are in parentheses. by Bhattacharyya & Nityananda (2008), we obtained phase- b Roy et al. (2012). connected time-of-arrivals (TOAs) from TEMPO,16 using the c Cordes & Lazio (2002). JPL DE405 solar system ephemeris (Standish 2004). The d Phase averaged. binary timing model used is ELL1 (Lange et al. 2001), since J1544 + 4937 is in a very low eccentricity system. This MSP is in a very compact binary with an orbital period of 2.9 hr. We One of the targets was Fermi J1544.2+4941, a γ -ray source derive a minimum companion mass (for 90◦ orbital inclination) from an unpublished internal source list created by the LAT of 0.017 M using the Keplerian mass function, assuming a Collaboration using 18 months of data in preparation for the pulsarmassof1.4M. J1544 + 4937 is eclipsed for about 13% 2FGL catalog (Nolan et al. 2012). The source location (J2000) of the orbit at 322 MHz (Section 6). The best-fit timing model from that analysis (and used for our telescope pointing) was ◦ ◦ (MJD 55680.927–56332.90) is obtained excluding the TOAs R.A. = 236.074, decl. = 49.695, with a 95% confidence error − around the eclipse phase (0.05 0.35). We achieved a post-fit circle of radius 9.5. This source is very weak, with a likelihood rms timing residual of 6.9 μs from 652 days of timing (Figure 1). test statistic (TS; Mattox et al. 1996) of 26.2 in the 18 month There are still unmodeled residuals, which can be partially analysis, and did not make the significance cut to be included in absorbed by proper motion fit. However, the inclusion of proper the 2FGL catalog itself. motion reduces the LAT detection significance, indicating that We processed the data on an IUCAA HPC cluster with more timing data are required to improve the model. We estimate Fourier-based acceleration search methods using PRESTO a precise DM equal to 23.2258(11) pc cm−3 by a timing (Ransom et al. 2002). We investigated trial dispersion measures −3 −3 fit using 322 and 607 MHz TOAs from non-eclipsing binary (DMs) ranging from 0 pc cm up to 350 pc cm . A linear drift phases. Ephemeris, position, and derived parameters are listed of up to 200 Fourier-frequency bins for the highest summed har- in Table 1. monic was allowed. The powerline, 50 Hz, and its subsequent harmonics were excised. Using parameters of 32 MHz band- 16 http://tempo.sourceforge.net

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Scaling attractors in interacting teleparallel dark energy

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Please note that terms and conditions apply. JCAP07(2013)044 hysics P le ic t ) an arbitrary function of ar φ ( f doi:10.1088/1475-7516/2013/07/044 , with strop f √ A / ,φ f ≡ α osmology and and osmology C . It is shown that in this case the existence of scaling attractors is possible, 1305.0474 φ modified gravity, dark energy theory It has been proposed recently the existence of a non-minimal coupling between rnal of rnal ou An IOP and SISSA journal An IOP and 2013 IOP Publishing Ltd and Sissa Medialab srl Instituto de Estadual F´ısicaTeórica,UNESP-Univ Paulista Caixa Postal 70532-2, 01156-970 SãoPaulo, Brazil E-mail: [email protected] c G. Otalora Received May 3, 2013 Revised June 23, 2013 Accepted July 6, 2013 Published July 29, 2013 Abstract. Scaling attractors in interacting teleparallel dark energy a canonical scalar fieldmotivated (quintessence) by and similar gravity constructionsthe in in model, the the known framework contextattractor as of of has teleparallel teleparallel General been gravity, dark Relativity. found.ruled energy, by The has Here a dynamics we been dynamically consider of changing further a coefficient developed, model but in no which scaling the non-minimal coupling is

J the scalar ﬁeld which means that thethe universe initial will conditions. eventuallyalleviated without enter As ﬁne-tunings. these a scaling consequence, attractors, the regardless cosmological of Keywords: coincidence problem could be ArXiv ePrint: Home Search Collections Journals About Contact us My IOPscience

Isocurvature and curvaton perturbations with red power spectrum and large hemispherical asymmetry

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Please note that terms and conditions apply. JCAP07(2013)043 5 when . 0 can easily | > hysics | Φ and curvaton | P Φ | / | y of isocurvature Φ le | the SM, cosmological ic ation of t pherical asymmetry as a olving along the tachyonic ar s, University of Lancaster, hemispherical asymmetry at erved by WMAP and Planck. and the tachyonic mass term 10.1088/1475-7516/2013/07/043 l Physics, ependence requirements on the ith a large hemispherical asym- N doi: d power spectrum. Dark matter ings of the beginning of tachyonic nitially sub-horizon quantum ﬂuc- strop A across the horizon, with ∆ | Φ | osmology and and osmology C is in the range 0.1-1. The spectral index of the isocurvature c physics of the early universe, cosmology of theories beyond We calculate the power spectrum and hemispherical asymmetr

[email protected] rnal of rnal ou . We find that a large hemispherical asymmetry due to the modul An IOP and SISSA journal An IOP and 2 2013 IOP Publishing Ltd and Sissa Medialab srl Lancaster-Manchester-Sheffield Consortium for Fundamenta Cosmology and Astroparticle PhysicsLancaster Group, LA1 Dept. 4YB, of U.K. Physic E-mail: c and curvaton perturbations duepart to of a its complex potential.tuations, field Using Φ we a which compute semi-classical isfunction the evolution ev of power of the spectrum, i number mean of field e-foldings of and tachyonic hemis growth ∆ cH Received May 21, 2013 Accepted June 24, 2013 Published July 29, 2013 Abstract.

John McDonald J Isocurvature and curvaton perturbations with red power spectrum and large hemispherical asymmetry be generated via the spatial modulation of the observed Universe exits theevolution horizon and within 10-40 e-fold

perturbations is generallyisocurvature negative, perturbations corresponding due to tometry a may an re be axion-like able curvaton toIn explain w this the hemispherical case, asymmetrysmall obs the scales, red which spectrumasymmetry should can from make additionally quasar it number suppress counts. easier the to satisfy scale-d Keywords: perturbation theory, CMBR theory The Astrophysical Journal, 773:93 (20pp), 2013 August 20 doi:10.1088/0004-637X/773/2/93 C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

PATHWAYS OF LARGE-SCALE MAGNETIC COUPLINGS BETWEEN SOLAR CORONAL EVENTS∗

Carolus J. Schrijver1, Alan M. Title1, Anthony R. Yeates2, and Marc L. DeRosa1 1 Lockheed Martin Advanced Technology Center, 3251 Hanover Street, Palo Alto, CA 94304, USA 2 Department of Mathematical Sciences, Durham University, Science Laboratories, South Road, Durham DH1 3LE, UK Received 2013 March 26; accepted 2013 May 3; published 2013 July 29

ABSTRACT The high-cadence, comprehensive view of the solar corona by SDO/AIA shows many events that are widely separated in space while occurring close together in time. In some cases, sets of coronal events are evidently causally related, while in many other instances indirect evidence can be found. We present case studies to highlight a variety of coupling processes involved in coronal events. We find that physical linkages between events do occur, but concur with earlier studies that these couplings appear to be crucial to understanding the initiation of major eruptive or explosive phenomena relatively infrequently. We note that the post-eruption reconfiguration timescale of the large-scale corona, estimated from the extreme-ultraviolet afterglow, is on average longer than the mean time between coronal mass ejections (CMEs), so that many CMEs originate from a corona that is still adjusting from a previous event. We argue that the coronal field is intrinsically global: current systems build up over days to months, the relaxation after eruptions continues over many hours, and evolving connections easily span much of a hemisphere. This needs to be reflected in our modeling of the connections from the solar surface into the heliosphere to properly model the solar wind, its perturbations, and the generation and propagation of solar energetic particles. However, the large-scale field cannot be constructed reliably by currently available observational resources. We assess the potential of high-quality observations from beyond Earth’s perspective and advanced global modeling to understand the couplings between coronal events in the context of CMEs and solar energetic particle events. Key words: magnetic fields – Sun: corona – Sun: coronal mass ejections (CMEs) – Sun: flares Online-only material: color figures

1. INTRODUCTION an adequate distinguishing criterion for inferring causal links between events on the Sun. The Atmospheric Imaging Assembly (AIA; Lemen et al. In order to assess the importance of causal linkages in the 2012)ontheSolar Dynamics Observatory (SDO; Pesnell et al. triggering of near-synchronous events, we first must identify 2012) provides continuous full-disk observations of the solar and classify the types of pathways that may connect them. To chromosphere and corona in seven extreme-ultraviolet (EUV) that end, we have reviewed many near-synchronous events in channels, spanning a temperature range from some 20,000 K to SDO/AIA observations, and here present a selection of those in excess of 20 MK (Boerner et al. 2012). The 12 s cadence of to illustrate the causal linkages. With the available present-day the image stream with 4096 × 4096 pixel images at 0.6 arcsec observations, we can see that many of the events observed with pixel−1 provides unprecedented views of the various phenomena SDO/AIA either reveal direct magnetic connections between that occur within the evolving solar outer atmosphere. near-synchronously flaring or erupting regions, while others are With the full-disk coverage and the multitude of events highly suggestive of it. occurring on the Sun at any given time, it is not surprising The idea of causal linkage between flaring in different regions that many events are seen to occur synchronously or near- goes back to Richardson (1936). Richardson (1951, page 171), synchronously. In many cases, this (near-)synchronicity is a following upon his initial report, noted that “the formal statistical matter of chance. During the 976 day time interval from the start results as well as the visual impression conveyed from inspection of the SDO prime mission (2010 May 1) through the last of the of the [Ca ii K] photographs suggest that some form of coupling searched dates for this study (2012 December 31) the NOAA/ may exist between widely separated spot groups,” but with the GOES logs contained 2881 flares of class C1 or above, or on data at hand concluded that “the question must still be regarded average 3.0 flares per day. The study by Robbrecht et al. (2009) as open.” Since then, multiple studies were published on the counts typically some four coronal mass ejections (CMEs) for an possible linkages not only between flares but also between average day. For a causal link via an Alfvenic´ signal, we can take flares and filament activations and even between sets of filament a typical time delay Δτ for a signal to travel a distance d around activations; Dodson & Hedeman (1966), who list many of these the solar circumference with a characteristic coronal Alfven´ early studies, differentiate the linkage between relatively distant ◦ speed vA. For events separated by, e.g., 90 , d = 2πR/4, and regions from another form of linkage, namely that of multiple −1 with vA = 500 km s , we find that Δτ = 2200 s, or about flares occurring sequentially in the same region, which includes 1/40th of a day. With three flares and two CMEs a day on the subset of homologous flares (with a much larger scale the Earth-facing hemisphere, each of which can last for up to equivalent in repeating pseudo-streamer blowouts, e.g., Lynch multiple hours, it is not surprising that many events in the SDO/ & Edmondson 2013). AIA data are seen to overlap in time (quantitative estimates Among the possible scenarios for linkage between solar are provided in Section 5). Overlap or proximity in time is not impulsive events that were proposed already early on were fast- moving energetic particles and also shock waves traveling at up to 2000 km s−1. Evidence for the existence and possible role ∗ Supporting online materials can be found at http://www.lmsal.com/forecast/STYD.html. of these was found at radio and microwave wavelengths (e.g.,

1 The Astrophysical Journal, 773:93 (20pp), 2013 August 20 Schrijver et al. Wild 1969; Feix 1970). Thermal energy input conducted via discuss evidence for four fundamentally distinct plausible causal direct field connections was also proposed as a possible causal pathways: (evolving) direct magnetic connections, waves, or agent in sympathetic activity (e.g., Changxi et al. 2000). propagating disturbances, distortions of and reconnection with Larger data bases, some supported by space-based observa- overlying field by the eruption of one or more flux ropes tions of coronal connections, led to statistical studies in the elsewhere in the corona, and evolving indirect connections. context of connection patterns in the coronal magnetic field. One aim of this study is to present evidence for each of these Fritzova-Svestkova et al. (1976) and Pearce & Harrison (1990), pathways in at least one well-observed case. Our other aim is for example, following similar studies and conclusions by others to illustrate why it is proving difficult to assess the prevalence cited in their work, used those statistical methods on samples of sympathetic couplings in the causes of space weather and to of events to conclude that sympathetic flaring was at most a point toward future opportunities to alleviate the difficulties in weak phenomenon, although apparently significant for regions the study of this now nearly 80 yr old problem. in close proximity of each other (which those studies found to Section 2 describes the various data sets used in this study be closer than 30◦ and 35◦, respectively). and the selection criteria for initial review of candidate data sets. Suggestions of causal linkage between events are not limited Section 3 reviews the most illuminating cases that support causal to the flare-flare or flare-filament events, but also include linkage of solar events or clearly illustrate a particular difficulty couplings of filament eruptions. For example, Jiang et al. (2011) in establishing whether events are synchronous by chance or by discuss the possibility of two quiet-Sun filament eruptions interaction. In reviewing each of the events, we typically look occurring in the wake of an active-region eruption, with the at events over a full 24 hr period and sometimes up to several coupling agent being the field deformation by an eruption (often days. After reviewing Section 2 the reader may choose to jump with a signature “coronal dimming”) that is instrumental in to the final two sections, where we discuss our findings in a causing other field configurations that are either connected to summarizing Section 4 and in the concluding Section 5, before it or that lie contained within it to lose their stability. Another reviewing the detailed case studies in Section 3. To aid in the such pair of related quiet-Sun filament eruptions, along with review of the supporting images and movies, we created an other flaring activity, was discussed by Schrijver & Title (2011). online table (http://www.lmsal.com/forecast/STYD.html) and A magnetic configuration of three adjacent flux ropes was include direct links throughout the manuscript shown by a subsequently modeled by Tor¨ ok¨ et al. (2011), whose detailed superscript “S” followed by a number. MHD work illustrated how a stretched field by one eruption can destabilize one or more other flux-rope configurations nested 2. DATA within a common overlying field. Such linkages, particularly those involving fairly high field, are often not directly observable The primary data source for this study is the complete but can be inferred from models, as was done in the studies archive of SDO/AIA observations. We selected candidate events mentioned in this paragraph, and in the work by, for example, using several different approaches. First, we created 30 minute Yang et al. (2012) and Shen et al. (2012). summed images, downgraded to 64 × 64 pixels, in the 193 Å Despite the many studies cited above and referenced within channel, using the online 2 minute cadence synoptic data those, the phenomenon of sympathetic activity in the solar set at 1024 × 1024 pixels from 2010 May 1 through 2012 corona remains elusive. Yet, in the context of forecasting solar December 31 (http://jsoc.stanford.edu/data/aia/synoptic/). We activity, the influence of adjacent or distant regions on the loss chose the 193 Å channel because it reveals both flaring activity of stability of a given region needs to be understood. This is and the lower-energetic phenomena of eruptions from quiet- particularly important for the development and propagation of Sun regions. We then selected all events that exceeded a flare- CMEs and the resulting particle events and geomagnetic storms: like intensity and remapped their coordinates to an evolving given the frequency of eruptive events on the Sun, many CMEs synoptic map set, letting the signal fade over a two-week period are composite events, but understanding their makeup from to allow patterns to stand out clearly. All pronounced clusters of different events with either physical connections or chance events were subsequently reviewed in the daily summary movies coincidences is important both to the interpretation and the of AIA observations, in the process eliminating instrumental forecast of any heliospheric event under study. artifacts related to spacecraft rolls and off-points, data gaps, and The complete coverage of the Earth-facing hemisphere of calibration mode data. the solar surface and corona by SDO, supported by far-side As a next pre-selection criterion, we reviewed all M- and observations from the STEREO spacecraft (Kaiser et al. 2008), X-class flares in the AIA archive during the same period, and complemented by the STEREO and SOHO/LASCO (Brueckner once again reviewed all daily summary movies for those dates. et al. 1995) coronagraphs and full-sphere coronal field modeling We also collected candidate events during daily reviews of (such as work by Schrijver & DeRosa 2003; Yeates et al. 2008, AIA data as annotated into the Heliophysics Events Registry which we use in the present study) enable a comprehensive (Hurlburt et al. 2012). From the sample of some five dozen empirical assessment of the linkages that exist in the solar candidate events, we selected the 10 cases that were most com- corona that may play a role in sympathetic activity. In view of pelling by visual inspection in supporting causal connections. the mixture of positive and negative findings in the literature of For selected events, we also made and reviewed running- the significance of causal linkages between explosive or eruptive ratio movies in which the 211 Å, 193 Å, and 171 Å channels events on the Sun, the aim of this study is not to quantify the were combined. The frames in these image sequences were frequency of sympathetic couplings, but to assess the evidence created by first computing logarithmic differences for time- for, and to discuss examples of, any of the proposed causal averaged images (for 264 s averages of fixed-exposure frames pathways by which couplings may occur. taken at 24 s cadence and 2 minute offsets between successive In selecting the events discussed here, we reviewed much of frames to be differenced for these movies), and then combining the SDO/AIA data for the period from 2010 May 1 through 2012 these in sets of three into the rgb color planes of a movie December 31. Based on the review of those observations, we (clipping the scales at relative brightening or dimming to range

POPULATION III STARS AND REMNANTS IN HIGH-REDSHIFT GALAXIES

Hao Xu1,JohnH.Wise2, and Michael L. Norman1 1 Center for Astrophysics and Space Sciences, University of California, San Diego, 9500 Gilman Drive, La Jolla, CA 92093, USA; [email protected], [email protected] 2 Center for Relativistic Astrophysics, School of Physics, Georgia Institute of Technology, 837 State Street, Atlanta, GA 30332, USA; [email protected] Received 2013 May 5; accepted 2013 June 24; published 2013 July 29

ABSTRACT Recent simulations of Population III star formation have suggested that some fraction form in binary systems, in addition to having a characteristic mass of tens of solar masses. The deaths of metal-free stars result in the initial chemical enrichment of the universe and the production of the first stellar-mass black holes. Here we present a 9 cosmological adaptive mesh refinement simulation of an overdense region that forms a few 10 M dark matter halos and over 13,000 Population III stars by redshift 15. We find that most halos do not form Population III stars 7 until they reach Mvir ∼ 10 M because this biased region is quickly enriched from both Population III and galaxies, which also produce high levels of ultraviolet radiation that suppress H2 formation. Nevertheless, Population III −4 −1 −3 stars continue to form, albeit in more massive halos, at a rate of ∼10 M yr Mpc at redshift 15. The most 7 massive starless halo has a mass of 7 × 10 M, which could host massive black hole formation through the direct gaseous collapse scenario. We show that the multiplicity of the Population III remnants grows with halo 8 9 mass above 10 M, culminating in 50 remnants located in 10 M halos on average. This has implications that high-mass X-ray binaries and intermediate-mass black holes that originate from metal-free stars may be abundant in high-redshift galaxies. Key words: galaxies: star formation – hydrodynamics – methods: numerical – radiative transfer Online-only material: color figures

= 1. INTRODUCTION still form at the late epoch of z 6 in the underdense regions of the universe by post-processing of cosmological simulations The ﬁrst generation stars (Population III) form from with blast wave models. Muratov et al. (2012) also showed that 6 metal-free gas in dark matter halos with M ∼ 10 M and Pop III stars continue to form until z = 6 using direct cosmo- have a large characteristic mass (e.g., Abel et al. 2002;Bromm logical simulations. et al. 2002; O’Shea & Norman 2007;Turketal.2009;Greif In addition to the metal enrichment from Pop III stars, heating et al. 2012). Due to their high mass, they have short life- and ionizing effects from their radiation are crucial to modeling times (Schaerer 2002) and may go supernova, and enrich their early structure formation of the universe (Haiman et al. 2000). surrounding intergalactic medium (IGM). For these metal-free Lyman–Werner (LW) photons from Pop III stars photodissociate stars, Type II supernovae (SNe) happen for an initial mass be- H2 by the Solomon process and suppress the formation of Pop III tween 10 and 40 M, and much more energetic pair-instability stars in low-mass halos. The higher energy UV radiation from SNe (PISNe) might occur in stars between 140 and 260 M Pop III stars then affects the subsequent structure formation (Heger et al. 2003). Once the metallicity passes some critical through heating and ionizing the surrounding IGM (Machacek −6 metallicity, ∼10 Z if dust cooling is efﬁcient (Omukai et al. et al. 2001; Yoshida et al. 2003; Wise & Abel 2007b; O’Shea −3.5 2005; Schneider et al. 2006; Clark et al. 2008)or∼10 Z &Norman2008). Pop III more massive than 260 M or less otherwise (Bromm et al. 2001; Smith et al. 2009), the gas can massive than 140 M may directly collapse to form black holes cool rapidly and lower its Jeans mass. These metal-enriched (BHs; Heger et al. 2003). Accretion onto these massive Pop III (Population II) stars have a lower characteristic mass scale and BHs is a feasible way to form z>7 quasars (Johnson et al. most likely have an initial mass function (IMF) that resembles 2013) and is also an important source of X-ray radiation in high the present-day one. However at high redshift, heating from the redshifts. X-rays from the accretion onto Pop III BHs (Alvarez cosmic microwave background may limit the radiative cooling, et al. 2009) or from Pop III binaries (Turk et al. 2009;Stacy thus increasing the Jeans mass, resulting in an IMF that also et al. 2010;Stacy&Bromm2013) may preheat and pre-ionize favors massive star formation (Larson 2005; Smith et al. 2009). a large volume of the IGM (Ricotti & Ostriker 2004; Mesinger The transition from Pop III to Pop II star formation is solely et al. 2013). dependent on the metal enrichment from the Pop III SN rem- The impact of Pop III feedback on cosmic evolution is nants in the future star-forming halos. Metal enrichment involves dependent on the properties of their host halos, especially their complex interactions between SNe blast waves, the IGM, halo masses (Whalen et al. 2008; Muratov et al. 2012). The sizes of mergers, and cosmological accretion. This topic has been ex- the host halos determine the distance the metals from Pop III tensively studied with semi-analytic models (Scannapieco et al. SNe reach with their blast waves, and the escape fraction of 2003; Yoshida et al. 2004; Tumlinson 2006; Salvadori et al. their UV radiation. This makes a detailed study of the Pop III 2007; Komiya et al. 2010), post-processing of numerical sim- star and remnant distribution over a wide range of high-redshift ulations (Karlsson et al. 2008; Trenti et al. 2009), and direct galaxies necessary. numerical simulations (Tornatore et al. 2007; Ricotti et al. 2008; It is impossible to observe Pop III stars during their lifetime Maio et al. 2010; Wise et al. 2012b; Muratov et al. 2012). For at high redshift, but they might still possibly be detected directly example, Trenti et al. (2009) suggested that Pop III stars may by looking for their PISN explosions before their death. This

1 The Astrophysical Journal, 773:83 (8pp), 2013 August 20 Xu, Wise, & Norman idea has been studied and shown to be promising for both LSST stars are formed otherwise. We use the same star formation (Trentietal.2009) and James Webb Space Telescope (Hummel models and most of the parameters in Wise et al. (2012b), as et al. 2012; Whalen et al. 2013). Understanding the population well as feedback models. For the initial mass of Pop III stars, and distribution of Pop III stars and remnants in high-redshift we randomly sample from an IMF with a functional form galaxies is helpful in preparing the observation of these events. 1.6 In this paper, we focus on the formation of Pop III stars and − Mchar the population and multiplicity of Pop III stars and remnants f (log M)dM = M 1.3 exp − dM, (1) M in high-redshift galaxies. We have performed a simulation of a survey volume of over 100 comoving Mpc3 that includes a full primordial chemistry network, radiative cooling from which behaves as a Salpeter IMF above the characteristic mass, metal species, both Pop II and Pop III star formation and Mchar, but is exponentially cut off below that mass (Chabrier their radiative, thermal, mechanical, and chemical feedback. 2003; Clark et al. 2009). The only difference between the two The simulation runs on such a large volume from cosmological simulations is that we here use a characteristic mass of 40 M initial conditions, so that we can follow the formation and fate for the Pop III IMF, which is more in line with the latest results of Pop III stars in a statistically significant number of halos with of Pop III formation simulations (e.g., Turk et al. 2009;Greif a wide range of masses from a few million solar masses to one et al. 2012), instead of 100 M. Please see Sections 2.2 and 2.3 billion solar masses. We first describe our simulation model in of Wise et al. (2012b) for the details of the star formation and Section 2. Then in Section 3, we present our results of the Pop III stellar feedback used in the simulation, respectively. star and remnant distribution over early galaxies. We discuss the We show the evolution of the high-resolution region in findings and possible bias in our simulation in Section 4. Figure 1 at three redshifts, 25, 17.91, and 15. Here we show the large-scale structure in the inner 6.6 comoving Mpc in 2. SIMULATION SETUP the top row, and in the remaining rows, we focus on the most massive halo at each redshift by showing their density-weighted We perform the simulation using the adaptive mesh refine- projections of gas density, temperature, and metallicity. ment cosmological hydrodynamics code Enzo (O’Shea et al. 2004). Adaptive ray tracing (Wise & Abel 2011)isusedforthe 3. RESULTS radiation transfer of ionizing radiation, which is coupled to the hydrodynamics and chemistry in Enzo. The chemistry, cooling, We illustrate the evolution of the number of Pop III stars and and star formation and feedback models used in this simulation remnants, as well as their formation rate from the birth of the first are the same as in Wise et al. (2012b). PopIIIstartoz = 15 in the top panel of Figure 2. Massive Pop III We generate the initial conditions for the simulation us- stars have very short lifetimes (Schaerer 2002) and end their life ing Music (Hahn & Abel 2011)atz = 99 and use the cosmolog- by either directly collapsing to BHs or exploding as SNe (Heger ical parameters from the 7 yr Wilkinson Microwave Anisotropy et al. 2003), depending on their initial masses. More specifically, Probe ΛCDM+SZ+LENS best fit (Komatsu et al. 2011): ΩM = they die as Type II SNe if 11 M/M 40 and as PISNe if 0.266, ΩΛ = 0.734, Ωb = 0.0449, h = 0.71, σ8 = 0.81, and 140 M/M 260, where M is the initial stellar mass, n = 0.963. We use a comoving simulation box of (40 Mpc)3 or become BHs if their masses are not in these mass ranges. that has a 5123 root grid resolution and three levels of static Throughout the paper, we use “Pop III remnants” to refer to all nested grids. We first run a 5123 N-body only simulation to remains of dead Pop III stars, regardless of whether they become z = 6. Then we select the Lagrangian volume around two BHs or supernova remnants, and we use “Pop III” to refer to both 10 ∼3 × 10 M halos at z = 6, and re-initialize the simula- living stars and dead remnants. In the case that remnants have tion with three more nested grids to have an effective reso- negligible masses after SN events, their star masses are replaced lution of 40963 and an effective dark matter mass resolution with a very small mass proportional to their initial masses and the 4 of 2.9 × 10 M inside the highest nested grid with a comov- star particles are kept in the simulation to follow their remnant ing volume of 5.2 × 7.0 × 8.3Mpc3 (300 Mpc3). During the kinematic distribution. The first Pop III star forms at redshift course of the simulation, we allow a maximum refinement level z ∼ 29.7. Rates of Pop III star formation steadily increase from l = 12, resulting in a maximal resolution of 19 comoving pc. ∼10−8 to above 10−6 stars per year per comoving Mpc3.The The refinement criteria employed are also the same as in Wise Pop III star formation rate (SFR) shows some signs of saturation et al. (2012b). The refinements higher than the static nested at z ∼ 15 in this overdense region of the universe. At z = 15, the grids are only allowed in the Lagrangian volume of the two entire survey volume of 138 comoving Mpc3 contains 13,123 massive halos at z = 6, which contains only high-resolution Pop III stars and remnants and 7677 halos more massive than 3 6 particles. It has a comoving volume of 3.8 × 5.4 × 6.6Mpc 5 × 10 M. The number densities of Pop III and halos that host (∼138 Mpc3)atz = 15. This highly refined volume is also the Pop III are 95 and 55 per comoving Mpc3, respectively. survey volume of this study. At this time, the simulation has The total mass of Pop II stars and the SFR of Pop II and 8 a large number (∼1000) of halos with M>10 M, where Pop III are shown in the bottom panel of Figure 2. There is 8 new formation of Pop III stars declines rapidly, for statistical ∼3 × 10 M mass in Pop II stars at z = 15. The star formation analysis. We use results at this redshift for our current study. histories for both Pop II and Pop III stars are similar to those The simulation has 1.3 billion computational cells and required in Wise et al. (2012b), but shifted to higher redshifts. We will 9 more than 10 million CPU hours, and there are three >10 M study the details of Pop II star formation in this simulation in a halos in the refined regions at this time. We will continue this forthcoming paper. simulation to lower redshift for the study of the Pop II stars and We show the number of Pop III, all halos, and star-hosting high-redshift galaxies. halos as functions of halo mass inside the survey volume at Both Pop II and Pop III stars form in the simulation, and z = 17.91 and z = 15 in the left panels in Figure 3, while the we distinguish them by the total metallicity of the densest star- numbers of living Pop III stars are stacked over remnants in the forming cell. Pop III stars are formed if [Z/H] > −4, and Pop II same bin. In the right panels in Figure 3 are shown the fraction

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Triggers for displaced decays of long-lived neutral particles in the ATLAS detector

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Please note that terms and conditions apply. 2013 JINST 8 P07015 EDIALAB July 5, 2013 M July 29, 2013 : May 10, 2013 : : ISSA S CCEPTED ECEIVED A UBLISHED R P UBLISHINGFOR doi:10.1088/1748-0221/8/07/P07015 IOPP UBLISHEDBY P : Trigger concepts and systems (hardware and software); Online farms and online : A set of three dedicated triggers designed to detect long-lived neutral particles de- CERN 2013 for the benefit of the ATLAS collaboration, published under the terms [email protected] c of the Creative Commons Attribution 3.0 License by IOP Publishing Ltd and Sissa BSTRACT EYWORDS E-mail: A K Medialab srl. Anypublished further article’s distribution title, journal of citation this and work DOI. must maintain attribution to the author(s) and the caying throughout the ATLAS detector tothe a triggers pair for of selecting hadronicsimulated displaced jets events. is decays The described. as effect The a ofof efficiencies pile-up function of the interactions of trigger on the rate the on trigger decaydiscussed. instantaneous efficiencies position and luminosity are the during dependence presented the 2012 for data-taking period at the LHC are filtering; Trigger algorithms The ATLAS collaboration Triggers for displaced decays of long-livedparticles neutral in the ATLAS detector The Astrophysical Journal, 773:90 (14pp), 2013 August 20 doi:10.1088/0004-637X/773/2/90 C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

STELLAR VELOCITY DISPERSION MEASUREMENTS IN HIGH-LUMINOSITY QUASAR HOSTS AND IMPLICATIONS FOR THE AGN BLACK HOLE MASS SCALE

C. J. Grier1, P. Martini1,2,8, L. C. Watson3, B. M. Peterson1,2,M.C.Bentz4, K. M. Dasyra5, M. Dietrich6, L. Ferrarese7, R. W. Pogge1,2, and Y. Zu1 1 Department of Astronomy, The Ohio State University, 140 W 18th Avenue, Columbus, OH 43210, USA 2 Center for Cosmology and AstroParticle Physics, The Ohio State University, Columbus, OH 43210, USA 3 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA 4 Department of Physics and Astronomy, Georgia State University, Atlanta, GA 30303, USA 5 Observatoire de Paris, LERMA (CNRS:UMR8112), 61 Avenue de l’Observatoire, F-75014, Paris, France 6 Department of Physics and Astronomy, Ohio University, Athens, OH 45601, USA 7 Herzberg Institute of Astrophysics, National Research Council of Canada, 5071 West Saanich Road, Victoria BV V9E 2E7, Canada Received 2013 May 2; accepted 2013 June 13; published 2013 July 29

ABSTRACT We present new stellar velocity dispersion measurements for four luminous quasars with the Near-Infrared Integral Field Spectrometer instrument and the ALTAIR laser guide star adaptive optics system on the Gemini North 8 m telescope. Stellar velocity dispersion measurements and measurements of the supermassive black hole (BH) masses in luminous quasars are necessary to investigate the coevolution of BHs and galaxies, trace the details of accretion, and probe the nature of feedback. We ﬁnd that higher-luminosity quasars with higher-mass BHs are not offset with respect to the MBH–σ∗ relation exhibited by lower-luminosity active galactic nuclei (AGNs) with lower-mass BHs, nor do we see correlations with galaxy morphology. As part of this analysis, we have recalculated the virial products for the entire sample of reverberation-mapped AGNs and used these data to redetermine the mean virial factor f that places the reverberation data on the quiescent MBH–σ∗ relation. With our updated measurements and new additions to the AGN sample, we obtain f =4.31 ± 1.05, which is slightly lower than, but consistent with, most previous determinations. Key words: galaxies: active – galaxies: kinematics and dynamics – galaxies: nuclei – quasars: individual (PG1411+442, PG1617+175, Mrk 509, PG 2130+099) Online-only material: color ﬁgures

1. INTRODUCTION Di Matteo et al. 2005, 2008). The MBH–σ∗ relationship can be used to infer MBH in large samples of galaxies. This allows for Over the past couple of decades, both observational and the exploration of the BH mass function on much larger scales analytical work have suggested a physical connection between (e.g., Yu & Tremaine 2002) and thus helps investigate the role the formation and growth of galaxies and the growth of their of BHs in galaxy formation and evolution processes. central black holes (BHs). For example, the comoving emissivity Direct MBH measurements are made with stellar kinematics of active galactic nuclei (AGNs) and the cosmic star formation and gas dynamics, although these methods require good spatial rate have both similarly declined since z ∼ 1(Boyleetal.1998; resolution and are presently only feasible for nearby galaxies. Franceschini et al. 1999; Merloni et al. 2004; Silverman et al. AGNs, however, offer the most robust tracer of the evolution 2008), which might imply a link between star formation and of the BH population over much of the history of the universe. AGN activity. In addition, luminous AGNs are more often found Under the assumption that the motion of the gas in the broad in massive early-type galaxies with young stellar populations line region (BLR) of AGNs is dominated by the gravitational (e.g., Sanders et al. 1988; Kauffmann et al. 2003; Veilleux et al. influence of the BH, one can use the virial relation MBH = Δ 2 2009). Further support of a BH–galaxy connection comes in the (fRBLR V )/G to obtain MBH, where RBLR is the average radius form of a number of correlations between properties of the host of the emitting gas in the BLR, usually either determined galaxies and the masses of their central BHs. A key relationship with reverberation mapping (RM; e.g., Peterson et al. 2004)or is between BH mass (MBH) and bulge stellar velocity dispersion estimated with the radius–luminosity relation (e.g., Bentz et al. (σ∗), observed in both quiescent (Ferrarese & Merritt 2000; 2009a, 2013), ΔV is the velocity dispersion of the gas, deduced Gebhardt et al. 2000a; Tremaine et al. 2002;Gultekin¨ et al. from the width of the emission line, and f is a dimensionless 2009; McConnell et al. 2011; McConnell & Ma 2013) and factor that accounts for the unknown geometry and orientation active galaxies (Gebhardt et al. 2000b; Ferrarese et al. 2001; of the BLR and may be different for each AGN. Nelson et al. 2004; Onken et al. 2004; Dasyra et al. 2007; With current technology, we are unable to directly observe Woo et al. 2010; Graham et al. 2011;Parketal.2012). This the structure of the BLR, as it is unresolvable even with the relation was first predicted by Silk & Rees (1998) and Fabian largest telescopes, so the true value of f for each object is (1999) and has been explained by various analytic models (e.g., unknown. This has contributed significantly to the uncertainties King 2003, 2005; Murray et al. 2005)aswellasrecoveredin in MBH measurements using BLR emission lines. Recent RM numerical simulations of evolving and interacting galaxies (e.g., efforts have begun to reveal more information about the actual structure of the BLR and the value of f in some objects (e.g., 8 Visiting Astronomer, North American ALMA Science Center and Bentz et al. 2010;Breweretal.2011; Pancoast et al. 2012;Grier University of Virginia, Charlottesville, VA 22903, USA. et al. 2013). However, limited data for most AGNs requires the

1 The Astrophysical Journal, 773:90 (14pp), 2013 August 20 Grier et al. use of an average virial factor f to estimate MBH. Currently, Table 1 f is calculated with the assumption that AGNs follow the Quasar Properties same MBH–σ∗ relation as quiescent galaxies (Onken et al. 2004; Galaxy R.A. Decl. z Woo et al. 2010, 2013; Graham et al. 2011;Parketal.2012). (J2000) (J2000) (NED) Most estimates of f are somewhat larger than ∼5; Onken PG 0026+129 00 29 13.6 +13 16 03 0.142 et al. (2004) find f =5.5 ± 1.8, Woo et al. (2010) find = ± PG 0052+251 00 54 25.1 +25 25 38 0.154 f 5.2 1.2, and more recently, analysis by Park et al. PG 1226+023 12 29 06.7 +02 03 09 0.158 (2012) and Woo et al. (2013) both yield f =5.1. Graham PG 1411+442 14 13 48.3 +44 00 14 0.089 = +0.7 et al. (2011) obtain a slightly lower value, f 3.8−0.6. PG 1617+175 16 20 11.3 +17 24 28 0.112 The difference between slopes and virial factors among PG 1700+518 17 01 24.8 +51 49 20 0.292 − studies using similar regression methods (whether MBH is Mrk 509 20 44 09.7 10 43 25 0.034 considered the independent or dependent variable) arise when PG 2130+099 21 32 27.8 +10 08 19 0.063 different galaxy samples are used to determine these two quantities, which may suggest a morphological dependence or Integral Field Spectrometer (NIFS) combined with the Gemini selection bias in the relation. In fact, recent studies do report North laser guide star AO system, ALTAIR. Watson et al. a morphological dependence in the quiescent MBH–σ∗ relation, (2008) used NIFS+ALTAIR to measure σ∗ for PG 1426+015 such that there are systematic differences in the relation for with much higher precision than previous measurements for early-type (higher-mass) and late-type (lower-mass) galaxies high-luminosity quasars. This success prompted us to undertake (e.g., Greene et al. 2010; McConnell & Ma 2013). Others additional observations of quasars at the high-mass end of the have found that barred galaxies lie systematically below the MBH–σ∗ relation. In this paper we present the results of our NIFS MBH–σ∗ relation of normal unbarred galaxies (e.g., Graham observations of eight high-luminosity quasars. We successfully 2008a, 2008b; Graham & Li 2009), and still others have found measured σ∗ in four objects and use these results to improve deviations in both the slope and intercept for galaxies hosting the population of the MBH–σ∗ relation at the high-mass end. pseudobulges (e.g., Hu 2008; Gadotti & Kauffmann 2009; We also recalculate virial products for the entire AGN sample Kormendy et al. 2011). The idea of a non-universal MBH–σ∗ with updated time lag measurements to re-derive f , calibrate relation has been supported by theoretical work as well (e.g., BH masses in AGNs, and reexamine the AGN MBH–σ∗ relation. King 2010; Zubovas & King 2012), which has also suggested In this work we adopt a cosmological model of Ωm = 0.3, −1 −1 that the relation may depend on environment. ΩΛ = 0.70, and H0 = 70 km s Mpc . Morphological deviations from a single MBH–σ∗ relation have also been claimed in AGNs (e.g., Graham & Li 2009; Mathur 2. OBSERVATIONS AND DATA ANALYSIS et al. 2012), and there has been some question as to whether or 2.1. NIFS/ALTAIR Observations not objects at the high-mass/high-σ∗ end of the relation follow a different slope (e.g., Dasyra et al. 2007;Watsonetal.2008). Observations of eight quasars were carried out at the Gemini 8 For example, four out of the six objects with MBH above 10 M North telescope in 2008 and 2010 under the programs included in the study of Watson et al. (2008) lie significantly GN-2008B-Q-28, GN-2010A-Q-11, and GN-2010B-Q-24. We above the relation. The appearance of outliers could be due chose our sample from the database of objects with RM- to systematic errors in σ∗ or MBH measurements, or simply a based BH mass measurements from Peterson et al. (2004) with 8 fluke due to small number statistics. Alternatively, Lauer et al. MBH > 10 M. Basic information on our targets is given in (2007) suggest that offsets at the high-mass end may be due to a Table 1. We used NIFS in conjunction with the ALTAIR laser selection bias. Specifically, when a sample is selected based on guide star AO system to carry out our observations. NIFS has a AGN properties, one is more likely to find a high-mass BH in 3 × 3 field of view that is divided into 29 individual spectro- a lower-mass galaxy (based on a BH–host galaxy correlation) scopic slices, with a spectral resolution R = λ/Δλ ≈ 5290 in because high-mass galaxies are rare and there is intrinsic scatter both the H and K bands. With the AO correction, NIFS yields a in BH–host galaxy correlations. spatial resolution on the order of 0.1. There are several strong An important step in evaluating the MBH–σ∗ relation and stellar absorption lines that fall within the wavelength range of any possible deviations from it is to obtain secure σ∗ and the H band, which has a central wavelength of 1.65 μm and MBH measurements in AGNs that sample the entire mass range covers from about 1.49 μmto1.80μm, so we observed seven of the relation. While the high-mass end of the quiescent of our targets with the H band filter. We list the most prominent 9 MBH–σ∗ relation is relatively well-populated to beyond 10 M stellar absorption features in this wavelength region in Table 2. (McConnell & Ma 2013), the current sample of AGNs used We observed our eighth object, PG 1700+518, in the K band due to calculate f still contains just three or four objects with to its higher redshift. The K filter on NIFS covers from about 8 MBH above 10 M (Graham et al. 2011;Parketal.2012;Woo 1.99 μmto2.40μm. et al. 2013). More measurements for luminous AGNs are needed We estimated the integration time for each object with HST/ to measure the high-mass end of the AGN MBH–σ∗ relation. ACS or WFPC2 images of the sources from Bentz et al. (2009a). However, accurate σ∗ measurements for high-luminosity AGNs To simulate the data we would obtain from NIFS, we measured are difficult to obtain because the AGN light overpowers the light the flux within a 3 × 3 aperture, except for a central circle from the host. Moreover, more luminous AGNs are relatively of diameter 0.2. We estimated the exposure time for each ob- scarce and thus typically found at large distances, so the host ject based on its brightness relative to PG 1426+015, for which galaxy has a small angular size and is easily lost in the glare Watson et al. (2008) obtained a host-galaxy signal-to-noise ra- of the AGN. It is only in the past few years that high-precision tio (S/N) ∼ 200 in about 2 hr of on-source integration. With measurements in very luminous objects have been obtained on both Poisson and background-limited trials, we estimated the account of the availability of adaptive optics (AO) and integral integration time required for each object to obtain a S/N ∼ field spectrographs (IFUs) such as Gemini North’s Near-Infrared 200. Table 3 gives details of the observations, most notably the

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Beam studies of novel THGEM-based potential sampling elements for Digital Hadron Calorimetry

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Please note that terms and conditions apply. 2013 JINST 8 P07017 d EDIALAB M July 29, 2013 May 19, 2013 June 26, 2013 June 20, 2013 : : : : for muons and a adout electron- ISSA 7 E. Oliveri, RS readout elec- S − a EVISED 10 ECEIVED R CCEPTED UBLISHED R A ∼ P detectors are presented. , THGEM, RETHGEM, ystems for particle and as- L. Moleri, . These results, together with 2 10.1088/1748-0221/8/07/P07017 c kHz cm and A. Breskin UBLISHING FOR doi: their suitability for economic pro- c efficiencies in the 98% range were y competitive compared to the other IOP P robabilities were tute of Science, 1. The resistive anode resulted in efficient dis- . C.D.R. Azevedo, 1 tribution of this b J.F.C.A. Veloso ∼ UBLISHED BY b P Creative Commons Attribution 3.0 ished article’s title, journal citation and DOI. H. Natal da Luz, a J.M.F. dos Santos, a configurations with a total thickness of 5–6 mm (excluding re pads inductively coupled through a resistive layer to APV-S 2 2 L. Arazi, 1 , 1cm by IOP Publishing Ltd and Sissa Medialab srl. Any further dis a 10cm : Calorimeters; Micropattern gaseous detectors (MSGC, GEM × : Beam studies of thin single- and double-stage THGEM-based CERN 2013, published under the terms of the × [email protected] A. Rubin, c License

a for pions in the double-stage conﬁguration, at rates of a few 6 Corresponding author. 1 − EYWORDS BSTRACT P.O. Box 26, Rehovot 76100, Israel Physics Department, FCTUC, University of Coimbra, 3004-516 Coimbra, Portugal I3N, Physics Department, University of Aveiro, 3810-193 Aveiro Portugal CERN, Meyrin, Switzerland E-mail: Dept. of Astrophysics and Particle Physics, Weizmann Insti c b a d K A Several 10 charge damping, with few-volt potential drops; discharge p 10 tronics, were investigated withrecorded muons with and an pions. average pad-multiplicity Detection of the robustness of THGEM electrodes againstduction spark over damage large and areas, make THGEM-basedtechnologies detectors considered highl for the SiD-DHCAL. ics), with 1 Beam studies of novel THGEM-based potential sampling elements for Digital Hadron Calorimetry S. Bressler, work must maintain attribution to the author(s) and the publ MHSP, MICROPIC, MICROMEGAS, InGrid, etc);troparticle physics; Large Gaseous detectors detector s M. Pitt, The Astrophysical Journal, 773:92 (21pp), 2013 August 20 doi:10.1088/0004-637X/773/2/92 C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

AN XMM-NEWTON SURVEY OF THE SOFT X-RAY BACKGROUND. III. THE GALACTIC HALO X-RAY EMISSION

David B. Henley and Robin L. Shelton Department of Physics and Astronomy, University of Georgia, Athens, GA 30602, USA; [email protected] Received 2013 February 20; accepted 2013 June 8; published 2013 July 29

ABSTRACT We present measurements of the Galactic halo’s X-ray emission for 110 XMM-Newton sight lines selected to minimize contamination from solar wind charge exchange emission. We detect emission from few million degree gas on ∼4/5 of our sight lines. The temperature is fairly uniform (median = 2.22 × 106 K, interquartile range = 0.63 × 106 K), while the emission measure and intrinsic 0.5–2.0 keV surface brightness vary by over an order of magnitude (∼(0.4–7) × 10−3 cm−6 pc and ∼(0.5–7) × 10−12 erg cm−2 s−1 deg−2, respectively, with median detections of 1.9×10−3 cm−6 pc and 1.5×10−12 erg cm−2 s−1 deg−2, respectively). The high-latitude sky contains a patchy distribution of few million degree gas. This gas exhibits a general increase in emission measure toward the inner Galaxy in the southern Galactic hemisphere. However, there is no tendency for our observed emission measures to decrease with increasing Galactic latitude, contrary to what is expected for a disk-like halo morphology. The measured temperatures, brightnesses, and spatial distributions of the gas can be used to place constraints on models for the dominant heating sources of the halo. We provide some discussion of such heating sources, but defer comparisons between the observations and detailed models to a later paper. Key words: Galaxy: halo – ISM: structure – X-rays: diffuse background – X-rays: ISM Online-only material: color ﬁgures

= ◦ 1. INTRODUCTION from archival XMM-Newton observations between l 120 and 240◦ (Henley & Shelton 2010, hereafter HS10). They compared Observations of the ∼0.1–1 keV diffuse soft X-ray back- the observed X-ray temperatures and emission measures of ground (SXRB; e.g., McCammon & Sanders 1990) show that the hot halo with the distributions expected from different ∼(1–3) × 106 K plasma is present in the halo of the Milky physical models. HSKJM’s analysis favored fountains of hot Way (Burrows & Mendenhall 1991; Wang & Yu 1995; Pietz gas (Joung & Mac Low 2006) as a major, possibly dominant, et al. 1998; Wang 1998; Snowden et al. 1998, 2000; Kuntz & contributor to the halo X-ray emission in the XMM-Newton Snowden 2000; Smith et al. 2007b; Galeazzi et al. 2007; Henley band over extraplanar SN remnants (Shelton 2006). However, & Shelton 2008; Lei et al. 2009; Yoshino et al. 2009; Gupta in the absence of X-ray surface brightness predictions from et al. 2009; Henley et al. 2010). The presence of this hot disk galaxy formation models, they were unable to rule out plasma is confirmed by the observation of zero-redshift O vii the possibility that an extended halo of accreted material also and O viii absorption lines in the X-ray spectra of active galac- contributed to the observed emission (Crain et al. 2010). tic nuclei (AGNs; Nicastro et al. 2002; Fang et al. 2003, 2006; Here, we expand upon HSKJM’s observational analysis, Rasmussen et al. 2003; McKernan et al. 2004; Williams et al. analyzing ∼4 times as many sight lines. Our observations 2005; Bregman & Lloyd-Davies 2007; Yao & Wang 2007;Yao are drawn from a new XMM-Newton SXRB survey which et al. 2008; Hagihara et al. 2010; Sakai et al. 2012; Gupta et al. spans the full range of Galactic longitudes (Henley & Shelton 2012). The extent and mass of this hot gas are disputed: Gupta 2012, hereafter HS12), and which supersedes the HS10 survey et al. (2012) argue that its extent is 100 kpc and that it con- from which the HSKJM sample was drawn. As in HSKJM, tains a significant fraction of the Galaxy’s baryonic mass (see our observations were selected because they should be less also Fang et al. 2013), while Wang & Yao (2012, and references affected by solar wind charge exchange (SWCX) emission therein) argue that the scale height of the hot gas is only a (Cravens 2000; Robertson & Cravens 2003; Koutroumpa et al. few kpc, in which case it would contribute a negligible amount 2006)—time-variable X-ray line emission which arises within to the Galaxy’s baryons (Fang et al. 2013). the solar system from charge exchange reactions between solar The origin of this hot halo gas is uncertain. Two main wind ions and neutral H and He (Cravens et al. 2001; Wargelin processes are thought to play a role in heating the halo. The et al. 2004; Snowden et al. 2004; Koutroumpa et al. 2007; first is supernova (SN) driven outflows from the Galactic disk Fujimoto et al. 2007; Kuntz & Snowden 2008; Henley & Shelton (e.g., Shapiro & Field 1976; Bregman 1980; Norman & Ikeuchi 2008; HS10; Carter & Sembay 2008; Carter et al. 2010, 2011; 1989; Joung & Mac Low 2006). In such an outflow, the material Ezoe et al. 2010, 2011). In a separate paper, we will use these may subsequently fall back to the disk in a so-called galactic observations to test models of the hot halo gas (D. B. Henley fountain. The second process is accretion of material from the et al., in preparation). intergalactic medium (e.g., Toft et al. 2002; Rasmussen et al. The remainder of this paper is organized as follows. In 2009; Crain et al. 2010). However, the relative importance of Section 2, we describe the observation selection and data these two processes is not well known. reduction. In Section 3, we describe our spectral analysis Henley et al. (2010, hereafter HSKJM) tested models of the method. We present the results in Section 4. We discuss and hot halo gas using a sample of 26 SXRB spectra extracted summarize our results in Sections 5 and 6, respectively.

1 The Astrophysical Journal, 773:92 (21pp), 2013 August 20 Henley & Shelton 2. OBSERVATIONS means of the halo temperatures and emission measures of the individual sight lines. We tabulate these mean values as the 2.1. Observation Selection results for sight line 103 in Columns 12 and 13 of Table 1.Sim- The observations that we analyze here are a subset of those ilarly, the Galactic coordinates for this sight line are the means analyzed by HS12, who extracted SXRB O vii and O viii in- of the longitudes and latitudes for the individual sight lines. The tensities from 1880 archival XMM-Newton observations spread subsequent analysis will use the mean results for sight line 103. across the sky. In order to minimize SWCX contamination, After grouping together observations of the same sight line, we apply various filters to the data (HSKJM). In particular, and combining the results from sight lines 103.1 through 103.27 to minimize contamination from geocoronal SWCX and near- as described above, our set of 163 observations yields 110 Earth heliospheric SWCX, we only use the portions of the measurements of the halo’s temperature and emission measure. XMM-Newton observations during which the solar wind proton The locations of our sight lines on the sky are shown in Figure 4, flux was low or moderate. If excising the periods of high solar below. wind proton flux from an XMM-Newton observation resulted in Note that our sample of observations includes 20 of the too little usable observation time, the observation was rejected 26 observations analyzed in HSKJM. Of the remaining six (see Section 2.4 of HS12). After this solar wind proton flux fil- observations, five (0200960101, 0303260201, 0303720201, tering, 1003 observations are usable (HS12, Table 2). We apply 0303720601, 0306370601) are excluded due to our using additional filters to these observations as follows. We minimize a later date to define the beginning of the solar minimum heliospheric SWCX contamination by using only observations phase. The sixth observation (0305290201) is not included toward high ecliptic latitudes (|β| > 20◦) taken during solar in HS12’s catalog, and so is not included here, as it exhibits minimum (after 00:00UT on 2005 June 11). As we are inter- strong residual soft proton contamination (see Section 3.5 ested in the Galactic halo, we use only observations toward high of HS12). Observations 0306060201 and 0306060301 were Galactic latitudes (|b| > 30◦), and exclude observations toward analyzed independently in HSKJM, but here they are grouped the Magellanic Clouds, the Eridanus Enhancement (Burrows together (sight line 43). et al. 1993; Snowden et al. 1995), and the Scorpius-Centaurus (Sco-Cen) superbubble (Egger & Aschenbach 1995). Note that 2.2. Data Reduction although we do not explicitly exclude the observations identified The data reduction is described in Section 2 of HS12 (see also as being SWCX-contaminated by Carter et al. (2011, Table A.1), Section 3 of HS10). Here, we give an overview of the process. none of these observations are in our final sample. The data reduction was carried out using the XMM-Newton The above criteria result in 163 observations being selected Extended Source Analysis Software3 (XMM-ESAS; Kuntz & from HS12’s original set of 1003. The observation IDs, names Snowden 2008; Snowden & Kuntz 2011), as included in version 2 of the original targets, and pointing directions for these 163 ob- 11.0.1 of the XMM-Newton Science Analysis Software4 (SAS). servations are shown in Columns 2–5 of Table 1 (Columns 6–9 Note that we re-extracted all the EPIC-MOS spectra from contain additional observation information (Section 2.2) and scratch for the current analysis, using a lower source removal Columns 10–15 contain the spectral fit results (Section 4)). If flux threshold than in HS12 (see below). the original target was a bright X-ray source, we excised it from Each observation was first processed with the SAS emchain the data, since our goal is to measure the diffuse SXRB emis- script to produce a calibrated events list for each exposure. Then, sion in each XMM-Newton field (see Section 2.2). Note that the XMM-ESAS mos-filter script was used to identify and these 163 observations represent fewer than 163 different sight ◦ excise periods within each exposure that were affected by soft lines. If a set of observations are separated by less than 0.1, proton flaring. As indicated above, periods of high solar wind we group them into a single sight line, and then fit our spectral proton flux (>2 × 108 cm−2 s−1) were also removed from the model (Section 3.1) to all the observations simultaneously. In data. The usable MOS1 and MOS2 exposure times that remain such cases, the observations for a given sight line are listed in after this filtering are shown in Columns 6 and 8 of Table 1, the table on and below the row containing the sight line number respectively. (e.g., the results for sight line 20 were obtained by simultane- Because our goal is to measure the diffuse Galactic halo emis- ously fitting to the spectra from observations 0400920201 and sion, we removed bright sources from the XMM-Newton data. 0400920101). As described in HS10 (Section 3.3) and HS12 (Section 2.2), Our set of 163 observations includes a cluster of 28 observa- ◦ ◦ we identified and removed bright and/or extended sources that tions near (l,b) ≈ (326 , −58 ). These observations represent would not be adequately removed by the automated source re- 27 different sight lines, which we have numbered 103.1–103.27 moval (described below). If the source to be removed was the (sight line 103.8 consists of two observations). In order to avoid original observation target, we centered the exclusion region on oversampling this region of the sky in our subsequent analysis, the target’s coordinates; otherwise, the exclusion region was po- we treat these 27 sight lines as a single sight line, whose halo sitioned by eye. In all cases we used circular exclusion regions. temperature and emission measure are found from the weighted We chose the radii of these regions by eye, although in some cases we used surface brightness profiles to aid us. As noted in 1 This date, taken from HS12, was estimated using sunspot data from the HS10 and HS12, we erred on the side of choosing larger exclu- National Geophysical Data Center (http://www.ngdc.noaa.gov/stp/SOLAR/). sion regions, at the expense of reducing the number of counts Note that this date is later than the one used in HSKJM, as HS12 defined an in the SXRB spectra. “Intermediate” phase of the solar cycle between solar maximum and solar minimum. We did not define an end date for the solar minimum phase, as the In general, we used the same source exclusion regions that sunspot data imply that this phase lasted at least until the most recent we used in HS10 and HS12. These were chosen from a visual observation in the HS12 catalog (carried out on 2009 November 3–4). inspection of broadband X-ray images, which had undergone 2 In general, the target names were obtained from the FITS file headers. If the target name was abbreviated or truncated, we attempted to determine the full name of the target from SIMBAD (http://simbad.u-strasbg.fr/simbad/). For a 3 http://heasarc.gsfc.nasa.gov/docs/xmm/xmmhp_xmmesas.html small number of targets, we were unable to determine the full name. 4 http://xmm.esac.esa.int/sas/

DRIVERSOFHi TURBULENCE IN DWARF GALAXIES

Adrienne M. Stilp1, Julianne J. Dalcanton1, Evan Skillman2, Steven R. Warren3, Jurgen¨ Ott4,andBarbel¨ Koribalski5 1 Department of Astronomy, University of Washington, Box 351580, Seattle, WA 98195, USA 2 Minnesota Institute for Astrophysics, University of Minnesota, 116 Church Street SE, Minneapolis, MN 55455, USA 3 Department of Astronomy, University of Maryland, CSS Building, Room 1024, Stadium Drive, College Park, MD 20742-2421, USA 4 National Radio Astronomy Observatory, P.O. Box O, 1003 Lopezville Road, Socorro, NM 87801, USA 5 Australia Telescope National Facility, CSIRO Astronomy and Space Science, P.O. Box 76, Epping, NSW 1710, Australia Received 2013 March 21; accepted 2013 June 8; published 2013 July 29

ABSTRACT Neutral hydrogen (H i) velocity dispersions are believed to be set by turbulence in the interstellar medium (ISM). Although turbulence is widely believed to be driven by star formation, recent studies have shown that this driving mechanism may not be dominant in regions of low star formation surface density (ΣSFR), such those as found in dwarf galaxies or the outer regions of spirals. We have generated average H i line profiles in a number of nearby dwarfs and low-mass spirals by co-adding H i spectra in subregions with either a common radius or ΣSFR. We find that the individual spatially resolved “superprofiles” are composed of a central narrow peak (∼5–15 km s−1) with higher velocity wings to either side, similar to their global counterparts as calculated for the galaxy as a whole. Under the assumption that the central peak reflects the H i turbulent velocity dispersion, we compare measures of H i kinematics determined from the superprofiles to local ISM properties, including surface mass densities and measures of star formation. The shape of the wings of the superprofiles do not show any correlation with local ISM properties, which indicates that they may be an intrinsic feature of H i line-of-sight spectra. On the other hand, the H i velocity dispersion is correlated most strongly with baryonic and H i surface mass density, which points toward a gravitational origin for turbulence, but it is unclear which, if any, gravitational instabilities are able to operate efficiently in these systems. Star formation energy is typically produced at a level sufficient to drive H i turbulent −4 −1 −2 motions at realistic coupling efficiencies in regimes where ΣSFR 10 M yr kpc , as is typically found in inner spiral disks. At low star formation intensities, on the other hand, star formation cannot supply enough energy to drive the observed turbulence, nor does it uniquely determine the turbulent velocity dispersion. Nevertheless, even at low intensity, star formation does appear to provide a lower threshold for H i velocity dispersions. We find a pronounced decrease in coupling efficiency with increasing ΣSFR, which would be consistent with a picture where star formation couples to the ISM with constant efficiency, but that less of that energy is found in the neutral phase at higher ΣSFR. We have examined a number of potential drivers of H i turbulence, including star formation, gravitational instabilities, the magneto-rotational instability, and accretion-driven turbulence, and found that, individually, none of these drivers is capable of driving the observed levels of turbulence in the low ΣSFR regime. We discuss possible solutions to this conundrum. Key words: galaxies: dwarf – galaxies: irregular – galaxies: ISM – galaxies: kinematics and dynamics – ISM: kinematics and dynamics Online-only material: color figures, figure sets

1. INTRODUCTION 2009). This relationship seems to hold in the central regions of spiral galaxies, but breaks down at large radii where SF The local velocity dispersion of neutral hydrogen (H i)pro- intensity drops dramatically while H i velocity dispersions vides a good tracer of the small-scale kinematics of the inter- remain relatively constant (e.g., Boulanger & Viallefond 1992; stellar medium (ISM) in disk galaxies. The velocity dispersions van Zee & Bryant 1999;Tamburroetal.2009). For some of H i clouds tend to range between 5 and 15 km s−1 across a massive spiral disks, H i velocity dispersions at large radii have wide range of galaxy and ISM properties (e.g., Dickey et al. been tentatively attributed to turbulence induced by non-stellar 1990; van Zee & Bryant 1999;Tamburroetal.2009). These sources, such as the magneto-rotational instability (MRI; e.g., velocity dispersions are thought to be turbulent in nature (e.g., Sellwood & Balbus 1999; Zhang et al. 2012). This outer disk Mac Low & Klessen 2004) because they are much greater than regime exhibits star formation rate (SFR) intensities and H i the velocity dispersions expected for the stable thermal temper- velocity dispersions similar to those found in dwarf galaxies, atures in the ISM (1 km s−1 and7kms−1 for the cold and warm which tend to have solid-body rotation curves (e.g., Oh et al. neutral phases; e.g., Wolfire et al. 1995). H i velocity dispersions 2011) and therefore lack the shear required for the MRI to also tend to either remain constant or decrease with increasing function efficiently. Additionally, H i velocity dispersions in the galaxy radius in spirals and large dwarfs (Dickey et al. 1990; outskirts of some dwarfs are not necessarily correlated with Boulanger & Viallefond 1992; Petric & Rupen 2007;Tamburro optical features or star forming regions (e.g., Hunter et al. 1999, et al. 2009). 2001). Studies of the relationship between SF and H i velocity The origin of turbulent H i velocity dispersions remains dispersions in dwarf galaxies may therefore help address the uncertain, but many studies attribute H i turbulence to star question of what provides the energy to drive H i turbulence, formation (SF) and resulting supernova (SN) explosions (e.g., particularly in regions where other proposed turbulence drivers MacLow&Klessen2004; Tamburro et al. 2009; Joung et al. are inefficient.

1 The Astrophysical Journal, 773:88 (21pp), 2013 August 20 Stilp et al. Stilp et al. (2013a, hereafter Paper I) presented a method 1. Instrumental angular resolution smaller than 200 pc, to to characterize the average H i kinematics in dwarf galaxies avoid artificially broadening H i line-of-sight spectra at by co-adding individual line-of-sight profiles after removal coarser resolution. of the rotational velocity for a single galaxy, as also used 2. Velocity resolution Δv 2.6kms−1, to resolve the width by Ianjamasimanana et al. (2012). These “superprofiles” were of the H i line-of-sight spectra. composed of a central peak with higher-velocity wings to either 3. Inclination i<70◦, to avoid broadening H i line-of-sight side. We interpreted the central peak of the superprofile as spectra with rotation. representative of the average H i turbulent kinematics, with the 4. No noticeable contamination from the Milky Way or a higher velocity wings representing anomalous motions such as companion, to ensure that detected H i belongs to each expanding H i holes or other bulk flows. Our conclusions in galaxy. Paper I were limited by the fact that the superprofiles were 5. More than 10 independent beams across the galaxy above a generated on global scales, whereas H i velocity dispersions are signal to noise threshold of S/N > 5, to allow for accurate known to vary across the disk. determination of the peak of H i line-of-sight spectra. In this paper, we extend the technique presented in Paper I 6. Available ancillary far ultraviolet (FUV, Galaxy Evolution to analyze subregions of these same galaxies. By extending Explorer) data, to uniformly measure SFRs. our analysis to carefully chosen subregions, we can essentially increase the dynamic range of various quantities, such as ΣSFR, In this paper, we apply additional selection criteria to the ΣH i, and ΣSFR/ΣH i, which were forced to galaxy-wide averages spatially resolved superprofiles to ensure that they are robust. in our earlier analysis. In contrast, many of the proposed drivers We have empirically found that superprofiles with fewer than of turbulence are local phenomena. This approach therefore five contributing independent beams are too noisy to accurately provides a more direct assessment of which parameters influence determine superprofile parameters, which eliminates four low- H i kinematics. mass galaxies (DDO 125, M81 DwB, NGC 4163, and GR We first compute superprofiles in radial annuli to facilitate 8) from the Paper I sample. Second, the superprofiles in the comparison with Tamburro et al. (2009), who found that SF other disks of some of the larger galaxies exhibit “clean bowls” does not provide enough energy to drive H i velocity dispersions that hinder accurate parameterization. These clean bowls are outside the optical radius r25 of spiral galaxies. However, radius due to missing short-spacings at the VLA, and can be present is not necessarily a good proxy for local ISM properties like ΣSFR as negative flux on either side of the central peak in the and ΣH i in the low-mass dwarfs in our sample. We therefore also superprofiles generated at large radii for some galaxies. We generate superprofiles in regions of constant ΣSFR to maximize eliminate these superprofiles from our analysis, and note that the sensitivity of the effects that SF may have on H i kinematics. they usually occur past 2r25. Our study is complementary to that of Tamburro et al. (2009), General properties of the final sample are given in Table 1. as we focus on lower-mass galaxies and isolate regions with Galaxies are listed in order of decreasing total baryonic mass similar SF surface density. (Mbaryon,tot). We list (1) the galaxy name; (2) the H i survey The layout of this chapter is as follows. In Section 2,we from which data were taken; (3–4) the position in J2000 discuss the data used to generate spatially resolved superpro- coordinates; (5) distance in Mpc; (6) inclination in degrees; (7) files and other galaxy properties. In Section 3,wegiveabrief total baryonic mass, Mbaryon,tot; (8) total H i mass, MH i,tot;(9) overview of the method used to generate the superprofiles, SFR as determined from FUV+24 μm emission; (10) the peak present the spatially resolved superprofiles, and address their ΣSFR included in this analysis for each galaxy; (11) the optical −1 robustness. In Section 4, we compare the superprofile parame- radius at a B-band surface brightness of 25 mag arcsec (r25); ters to galaxy physical properties. In Section 5, we then discuss and (12) de Vaucouleurs T-type. All references are given in the relevant correlations and compare SF energy to H i energy. Paper I, with the exception of the inclination for Sextans B. ◦ Finally, we summarize the conclusions in Section 6. The i = 52 value given in Paper I is a poor match to the H i morphology; we adopt i = 30◦, which is a much better match to the properties of the H i disk. 2. DATA 2.2. Converting Ancillary Data to Physical Properties We use H i data from the Very Large Array ACS Nearby Detailed information about the data used and methodology for Galaxy Survey Treasury Program (“VLA-ANGST”; Ott et al. deriving galaxy physical properties are given in Paper I. For this 2012) and The H i Nearby Galaxy Survey (“THINGS”; Walter study in particular, we focus on the SFR surface density (Σ ); et al. 2008). Following Paper I, we convolve these data to a SFR the SFR per available H i mass (Σ /Σ ); and the H i and common physical resolution of 200 pc to ensure that we are SFR H i baryonic surface densities (Σ , Σ ). Briefly, we calculate sampling ISM properties on the same physical scale for our H i baryon Σ using FUV and 24 μm data from the Local Volume Legacy entire sample. All spatially resolved ancillary data have also SFR (LVL) Survey (Dale et al. 2009) and the methodology outlined been convolved to this resolution to ensure a robust comparison in Leroy et al. (2008) and Paper I. We derive stellar surface between H i kinematics and other ISM properties. We assume mass density from LVL 3.6 μm data and apply the conversion distances as listed in Table 1 of Paper I, as compiled from Ott factor given in Leroy et al. (2008) and Paper I. Finally, we derive et al. (2012), Walter et al. (2008), and Dalcanton et al. (2009). baryonic surface mass density by combining the H i gas mass, including a factor of 1.36 correction for helium, with the stellar 2.1. Initial Sample Selection mass. All surface densities are inclination-corrected. For each superprofile subregion, we calculate ΣH i, Σbaryon, Σ We select our analysis sample for this paper from a subset of and SFR by taking the total MH i, Mbaryon, and SFR in that galaxies in Paper I. Briefly, the selection criteria used in Paper I subregion divided by its inclination-corrected area. We also Σ Σ are as follows. calculate SFR/ H i by measuring the total SFR in that subregion

ANALYTICAL LIGHT CURVE MODELS OF SUPERLUMINOUS SUPERNOVAE: χ 2-MINIMIZATION OF PARAMETER FITS

E. Chatzopoulos1, J. Craig Wheeler1,J.Vinko1,2, Z. L. Horvath2, and A. Nagy2 1 Department of Astronomy, University of Texas at Austin, Austin, TX, USA; [email protected] 2 Department of Optics and Quantum Electronics, University of Szeged, Hungary Received 2013 May 2; accepted 2013 June 21; published 2013 July 29

ABSTRACT We present fits of generalized semi-analytic supernova (SN) light curve (LC) models for a variety of power inputs including 56Ni and 56Co radioactive decay, magnetar spin-down, and forward and reverse shock heating due to supernova ejecta–circumstellar matter (CSM) interaction. We apply our models to the observed LCs of the H-rich superluminous supernovae (SLSN-II) SN 2006gy, SN 2006tf, SN 2008am, SN 2008es, CSS100217, the H-poor SLSN-I SN 2005ap, SCP06F6, SN 2007bi, SN 2010gx, and SN 2010kd, as well as to the interacting SN 2008iy and PTF 09uj. Our goal is to determine the dominant mechanism that powers the LCs of these extraordinary events and the physical conditions involved in each case. We also present a comparison of our semi-analytical results with recent results from numerical radiation hydrodynamics calculations in the particular case of SN 2006gy in order to explore the strengths and weaknesses of our models. We find that CS shock heating produced by ejecta–CSM interaction provides a better fit to the LCs of most of the events we examine. We discuss the possibility that collision of supernova ejecta with hydrogen-deficient CSM accounts for some of the hydrogen-deficient SLSNe (SLSN-I) and may be a plausible explanation for the explosion mechanism of SN 2007bi, the pair-instability supernova candidate. We characterize and discuss issues of parameter degeneracy. Key words: circumstellar matter – stars: evolution – stars: mass-loss – supernovae: general Online-only material: color figures

1. INTRODUCTION of Chevalier (1982) and Chevalier & Fransson (1994; see also Chevalier & Fransson 2001) to estimate the luminosity input The discovery of superluminous supernovae (SLSNe; from forward and reverse shocks depositing kinetic energy into Quimby et al. 2007, 2013; Smith et al. 2007; Gal-Yam 2012) the CSM and the SN ejecta, respectively. The CSM interaction imposed challenges to the widely used mechanism of 56Ni and plus 56Ni and 56Co radioactive decay LC model was successfully 56Co radioactive decay diffusion (Arnett 1980, 1982, hereafter compared in CWV12 to some radiation hydrodynamics numer- A80, A82) as the typical power input of many observed SN light ical LC models within uncertainties and was used to reproduce curves (LCs) that do not display prominent plateaus. Attempts the LC of the SN 2006gy. The availability of easily computed to fit the LCs of some SLSNe provided estimates for the mass analytical models allows the development of a χ 2-minimization of radioactive nickel, MNi, needed to power the peak luminosity fitting code that can be used to fit the observed LCs of SLSNe that were close to or far exceeded corresponding estimates for and other interesting transients. This fitting procedure allows us the total mass of the SN ejecta (Smith et al. 2007; Chatzopoulos to estimate the physical parameters involved and their uncer- et al. 2011; Chatzopoulos & Wheeler 2012;seeGal-Yam2012 tainties and to assess parameter degeneracies, a severe problem for a review). The striking variety in LC shapes, peak luminosi- with multi-parameter models, either analytic or numerical. The ties, durations, decline rates, and spectral evolution makes the resulting models give us a hint of which power input mechanism determination of a consistent physical model for the SLSNe even is most likely involved in these extraordinary events. This work more challenging. Radiation hydrodynamics simulations of in- serves as a sequel to the work of CWV12 and aims to apply fits teractions of SN ejecta with massive circumstellar matter (CSM) of the models presented there to all SLSNe for which LCs were shells of various power-law density profiles (Moriya et al. 2011, available when the work was done. The parameters derived from 2013; Ginzburg & Balberg 2012) provided important insights to those fits may be used as a starting point for more accurate but the dependence of the main features of the resulting numerical computationally expensive numerical simulations in the attempt LCs on the model parameters and were used to reproduce the to understand the physics involved in these rare explosions. observed LCs of some SLSNe (SN 2005ap, SN 2006gy, and We organize the paper as follows. In Section 2 we summarize SN 2010gx). the analytical LC models that were presented by CWV12 Chatzopoulos et al. (2012, hereafter CWV12) presented and used in this work to fit observed LCs. We also present generalized semi-analytical models for SN LCs that take into comparisons of our semi-analytic SN ejecta–CSM interaction account a variety of power inputs such as thermalized magne- model from CWV12 with numerical LC models of SN 2006gy. tar spin-down and forward and reverse shock heating due to In Section 3 we describe our observational sample of SLSNe and SN ejecta–CSM interaction with some contribution from 56Ni SNe IIn and in Section 4 the fitting method that was incorporated and 56Co radioactive decay. CWV12 considered cases where in our χ 2-minimization fitting code and present an analysis of the photosphere of the diffusion mass is either expanding ho- how it calculates uncertainties and parameter degeneracy related mologously or stationary within an optically thick CSM. Their to the large parameter space. We also present model fits to all formalism was largely based on that of A80 and A82 to in- events in our sample. Finally, in Section 5 we summarize our corporate an approximation for radiative diffusion and on that conclusions.

1 The Astrophysical Journal, 773:76 (21pp), 2013 August 10 Chatzopoulos et al. 2. SIMPLE MODELS FOR SLSNe LIGHT CURVES powered by such input is given by the following formula: x The analytical SN LC models that we use to fit observed 2E 2 2 x + w L(t) = p e−[x +wx] e[x +wx] dx, (2) SLSN LCs are presented in detail in CWV12. Here, we give t (1 + yx)2 a review of the models and of their physical assumptions. The p 0 derivation of those models was largely based on the methods where x = t/td and y = td /tp with td again being an “effective” discussed in A80 and A82 making the assumptions of homolo- diffusion time, Ep the initial magnetar rotational energy, and gous expansion for the SN ejecta, centrally located power input tp = the characteristic timescale for spin-down that depends source, radiation pressure being dominant, and separability of on the strength of the magnetic field. For a fiducial moment of the spatial and temporal behavior. In the generalized solutions inertia (1045 gcm2), the initial period of the magnetar in units 50 −1 0.5 presented in CWV12 we have relaxed the criterion for homolo- of 10 ms is given by P10 = (2 × 10 erg s /Ep) . The dipole gous expansion of the ejecta and also considered cases for large magnetic field of the magnetar can be estimated from P10 and initial radius as may be the case for the progenitors of some = 2 0.5 tp as B14 (1.3P10/tp,yr) , where B14 is the magnetic field SLSNe, as well as cases where the photosphere is stationary 14 in units of 10 G and tp,yr is the characteristic timescale for within an optically thick CSM envelope, as may be the case spin-down in units of years. Therefore, for the MAG model the for luminous interacting SNe IIn. The CSM interaction models fitting parameters are Ep, tp, R0, and td. We note that this model include bolometric LCs for both optically thick and optically assumes that the input from the pulsar is thermalized in the thin situations, as appropriate. The generalized solutions are ejecta. Simulations of this process show that the energy may not presented for a variety of power input mechanisms, including thermalize, but be ejected as magnetohydrodynamic (MHD) jets those that have been proposed in the past. (Bucciantini et al. 2006), thus compromising the mechanism as The first power input mechanism considered is the radioactive a model for SLSNe. 56 56 decay of Ni and Co (hereafter the RD model) that leads For both the RD and the MAG models, the SN ejecta mass, to the deposition of energetic gamma rays that are assumed Mej, is given by the following equation: to thermalize in the homologously expanding SN ejecta. As presented in CWV12, the generalized LC model in this case has 3 βc M = vt2, (3) the following form: ej 10 κ d

2M 2 where Mej is the mass of the SN ejecta, β is an integration = Ni −[x +2wx]{ − L(t) e (Ni Co) constant equal to about 13.8, and c is the speed of light. The td x value of Mej for a particular SN determined by its LC is 2 [x +2wx ] −td /tNix × (w + x )e e td dx uncertain because of the uncertainties associated with κ.For 0 the purposes of this work we will adopt the Thomson electron x [x2+2wx] scattering opacity for fully ionized solar metallicity material + Co (w + x )e (κ ∼ 0.33 cm2 g−1). We also adopt as a ﬁducial value for 0 = −1 − − −2 the expansion velocity v 10,000 km s for the estimates × e td /tCox t dx }· 1 − e At , (1) d presented in the tables. The uncertainty of Mej has an important effect on the criterion M

A NON-LTE ANALYSIS OF THE HOT SUBDWARF O STAR BD+28◦4211. I. THE UV SPECTRUM

M. Latour1, G. Fontaine1, P. Chayer2, and P. Brassard1 1 Departement´ de Physique, UniversitedeMontr´ eal,´ Succ. Centre-Ville, C.P. 6128, Montreal,´ QC H3C 3J7, Canada; [email protected] 2 Space Telescope Science Institute, 3700 San Martin Drive, Baltimore, MD 21218, USA Received 2013 April 18; accepted 2013 June 12; published 2013 July 29

ABSTRACT We present a detailed analysis of the UV spectrum of the calibration star BD+28◦4211 using high-quality spectra obtained with the Hubble Space Telescope and Far-Ultraviolet Spectroscopic Explorer satellites. To this aim, we compare quantitatively the observed data with model spectra obtained from state-of-the-art non-LTE metal line- blanketed model atmospheres and synthetic spectra calculated with TLUSTY and SYNSPEC. We thus determine in a self-consistent way the abundances of 11 elements with well-defined lines in the UV, namely those of C, N, O, F, Mg, Si, P, S, Ar, Fe, and Ni. The derived abundances range from about solar to 1/10 solar. We find that the overall quality of the derived spectral fits is very satisfying. Our spectral analysis can be used to constrain ◦ rather tightly the effective temperature of BD+28 4211 to a value of Teff = 82,000 ± 5000 K. We also estimate = +0.3 conservatively that its surface gravity falls in the range log g 6.2−0.1. Assuming that the Hipparcos measurement for BD+28◦4211 is fully reliable and that our model atmospheres are reasonably realistic, we can reconcile our spectroscopic constraints with the available parallax measurement only if the mass of BD+28◦4211 is significantly less than the canonical value of 0.5 M for a representative post-extended horizontal branch star. Key words: stars: abundances – stars: atmospheres – stars: fundamental parameters – stars: individual (BD+28◦4211) – subdwarfs Online-only material: color figures, extended figure

1. INTRODUCTION H/He ratio (but with no metals) for different effective temperatures and surface gravities. This method is now largely used ◦ BD+28 4211 is a hot subdwarf O (sdO) star whose bright- and its ability to give reliable fundamental parameters (Teff, ness, high effective temperature, and relatively simple spectrum log g, and sometimes also N(He)/N(H)) for white dwarfs and have made it a standard star in the optical domain as well as hot subdwarf B stars has already been demonstrated (Bergeron a calibration star for UV space missions such as IUE, Hub- et al. 1992; Saffer et al. 1994). However, this method is rather ble Space Telescope (HST), and FUSE. Its status of standard tricky in the case of BD+28◦4211 because, like others sdO stars star implies that some of its observational properties are very and white dwarfs that have high temperatures, the model spectra well known. For example, high precision UBVRI magnitudes used generally suffer from the so-called Balmer line problem. have been presented by Landolt & Uomoto (2007). In addi- That is, the observed lines cannot be simultaneously matched tion, its parallax measurement in the Hipparcos catalog places with a unique model spectrum. In other words, each line needs +14 ◦ the star at 92−11 pc. While studying BD+28 4211 as a spec- a model of different effective temperature in order to be well trophotometric standard, Massey & Gronwall (1990) found it reproduced. Usually, the lowest lines in the series (like Hα and to have a faint red companion at a separation of 2.8. Moreover, Hβ) need a lower temperature, while the highest ones are better BD+28◦4211 has been extensively observed in the UV range by reproduced at higher temperatures. In the case of BD+28◦4211, the missions mentioned above and there are highly valuable data to give the “extreme” values, Hα was best reproduced at Teff available on that star. Ultraviolet spectra are precious tools for 50,000 K and H at around 85,000 K (Napiwotzki 1993). The au- studying the chemical composition via the numerous metallic thor also found a log g value of 6.2 and a solar helium abundance lines present in this wavelength range. It could be thought that to be appropriate values for BD+28◦4211. His best estimate of with this privileged status, BD+28◦4211 would have been thor- the effective temperature of ∼82,000 K was subsequently con- oughly studied and its physical parameters would be accurately firmed by Dreizler & Werner (1993) who checked that parts known, but this is not exactly the case. In this connection, it of the IUE UV spectrum of BD+28◦4211, showing Fe vi and has to be mentioned that because of its high effective temper- Fe vii lines, are properly reproduced at around 82,000 K. This ature (around 80,000 K), the local thermodynamic equilibrium time, the NLTE model atmospheres they used included metals, (LTE) approximation is inappropriate for model atmospheres namely carbon, oxygen, nitrogen, and iron group elements, that intended to represent the star. Instead, the more sophisticated were grouped together into six model atoms (one for each ion- and realistic approach of non-LTE (NLTE) has to be used. ization stages between iii and viii). With similar models, and Given the physical and technical difficulties associated with on the basis of the same IUE spectrum, Haas et al. (1996)es- that approach, however, more efforts remain to be made along timated the abundance of iron to be about 10 times subsolar that avenue in order to characterize better the atmosphere of while nickel was found to be nearly solar. According to them, this star. oxygen and nitrogen also have abundances near the solar value. The first determination of the effective temperature of The improved UV spectra taken with Space Telescope Imaging BD+28◦4211, using quantitative spectroscopy, has been made Spectrograph (STIS) on board the HST allowed us to derive a by Napiwotzki (1993) who estimated it to be around 82,000 K. solar abundance of manganese, while a sole line of chromium He determined this value by comparing the Balmer lines of the indicated an abundance between two and four times the solar star with those of NTLE model atmospheres with a variable value (Ramspeck et al. 2003).

1 The Astrophysical Journal, 773:84 (15pp), 2013 August 20 Latour et al. The inclusion of metallic elements in NLTE model atmo- the abundance of these extra elements. This was done with the spheres, though costly in terms of computation time and com- MODION program4 which uses the TOPBASE data (Lanz et al. plexity, allows not only for more realistic models, but also per- 1996). This program allows the user to choose the explicit en- mits to solve, at least in part, the Balmer line problem (Werner ergy levels and also build superlevels that are included in the 1996). We will address this issue in more details in the sec- model atom. We thus constructed in this way model atoms for ond paper of this series (M. Latour et al. 2013, in preparation). the following ions : F iii with 9 levels and 5 superlevels, F iv The optical spectrum of BD+28◦4211 is rather featureless in with 11 levels and 5 superlevels, F v with 19 levels and 3 su- comparison to its rich UV spectrum. Except for the Balmer and perlevels, F vi with 12 levels and 5 superlevels, Mg iii with 37 He ii lines, nothing else is seen at medium resolution in the levels and 3 superlevels, Mg iv with 29 levels and 5 superlevels, optical range. Because of this, BD+28◦4211 was chosen to be Mg v with 18 levels and 2 superlevels, Ar iv with 39 levels, part of an investigation of diffuse interstellar bands in OB stars Ar v with 25 levels, Ar vi with 20 levels, and Ar vii with 18 given its uncomplicated spectrum. However, the high resolution levels. All transitions (bound-bound and bound-free) between HIRES spectra of the star obtained for this investigation at the these levels are thus considered when the ions are included in a Keck I telescope show a lot of narrow absorption lines as well model atmosphere. as a handful of emission lines (Herbig 1999). The sharpness of Since Werner (1996) underlined the importance of using Stark the absorption lines allowed to set an upper limit on the star’s profiles for CNO lines when modeling atmospheres of hot stars rotational velocity, v sin i 4kms−1, which is quite slow. such as BD+28◦4211, we inspected our different model atoms As for the high-quality FUSE spectrum of BD+28◦4211, it to check what kinds of profiles were used. We found that the has only been used to study interstellar abundances in the line strongest transitions (often resonance lines) of each ion are of sight of the star (Sonneborn et al. 2002). To our knowledge, treated with Stark profiles while the weaker ones are represented this data set has not been exploited so far to better characterize by Doppler profiles. We then examined a synthetic spectrum that the star. Although we have a general idea of its atmospheric could represent BD+28◦4211 and identified its most prominent chemical composition, no comprehensive or systematic studies lines and made sure they were described with a Stark profile in were made on that star. With high-quality data available from the corresponding atomic model. This way we added a classic HST and FUSE in particular, we felt we should exploit them in Stark profile to a few more lines of some elements, namely order to reexamine the chemical composition of BD+28◦4211 C iv,Niv,Oiv and O v,Siiv, and P v. A striking effect and also try to constrain better its effective temperature and that Werner (1996) noticed when including Stark profiles was surface gravity by studying the ionization equilibrium of some the disappearance of a high-temperature bump around log m metallic elements. Getting a portrait of the chemical compo- of −3gcm−2 which was present when using Doppler profiles sition of BD+28◦4211 should also be a first step in studying only (see his Figure 1). Since this bump is not seen either in our the optical spectrum of the star, with appropriate model atmo- models (see our Figure 1), we believe that our atomic data are spheres, for the Balmer line problem. appropriate for the study of BD+28◦4211 or other hot stars. Note In the second section of this paper we describe our model that, even without modifying the original model atoms used by atmospheres. This is followed by a short description of the Lanz & Hubeny (2003, 2007), our temperature structures do not observational material we used and of our abundance analysis present this bump (see Figure 4 of Latour et al. 2011). in Section 3. We then discuss our attempts at constraining The inclusion of our metallic elements must be done “step the effective temperature and the surface gravity by using the by step” if we want to ensure the convergence of our models. ionization equilibrium of metals in the UV range as well as the Too drastic changes in the physical parameters of the model parallax distance in Section 4. Finally, we present a discussion atmosphere used as input and the one we want to compute will and conclusion in Section 5. prevent the latter from converging. Thus, when constructing a grid of these line-blanketed models, we end up with a number 2. MODEL ATMOSPHERES of “subgrids” including only some of the elements mentioned above. In the case of BD+28◦4211, we built five “subgrids” in 2.1. Characteristics of Our Model Atmospheres order to end up with our final fully blanketed one. For example, We have developed the capacity to compute large grids of we have a grid including only C, N, and O in solar abundances, NLTE metal line-blanketed model atmospheres over reason- from which comes one of the models plotted in Figure 1. able timescales (days to weeks) with our parallel versions of In Figure 1, we show the temperature stratification for models = = TLUSTY and SYNSPEC that run on a dedicated cluster of with Teff 82,000 K, log g 6.2, and having a solar helium computers (currently containing 320 processors). Our setup is abundance culled from three of our grids. These estimates of described in more details in Latour et al. (2011) and has not the atmospheric parameters come from the work of Napiwotzki changed since, apart from the increase in the number of pro- (1993). The first model is a “classical” pure H+He NLTE model cessors we have available. Our final fully blanketed model at- that shows the well-known outwardly rise of temperature near mospheres for BD+28◦4211 include the following ions (besides the surface (dotted curve). The second one includes C, N, and those of H and He): C ii to C v,Nii to N vi,Oii to O vii,Siiii to O (solid curve), while the third one includes all the elements of Si v,Piv to P vi,Siii to S vii,Feiv to Fe viii, and Ni iii to Ni vii. our final model (see above) besides nickel (dashed curve). At The highest ionization stage of each element is taken as a one- this point, all our elements have a solar abundance (Grevesse level atom. More information on the model atoms we used can & Sauval 1998). Though we plotted only three models in the be found on TLUSTY’s Web site3 and in Lanz & Hubeny (2003, figure, we examined the ones (having the same parameters) 2007). Since our thorough examination of BD+28◦4211’s UV from our other grids and concluded that adding S, P, and Si to spectrum revealed also lines of argon, magnesium and fluorine, the C, N, O only induces a minor drop of the temperature in − we needed to construct additional model atoms in order to study the outer layers (log m< 2). When we add nickel to models

3 http://nova.astro.umd.edu/Tlusty2002/tlusty-frames-data.html 4 http://idlastro.gsfc.nasa.gov/ftp/contrib/varosi/modion/README

THE TWO-DIMENSIONAL PROJECTED SPATIAL DISTRIBUTION OF GLOBULAR CLUSTERS. I. METHOD AND APPLICATION TO NGC 4261

R. D’Abrusco1, G. Fabbiano1, J. Strader2, A. Zezas1,3,4,S.Mineo1, T. Fragos1,P.Bonﬁni3, B. Luo5, D.-W. Kim1, and A. King6 1 Harvard-Smithsonian Astrophysical Observatory, 60 Garden Street, Cambridge, MA 02138, USA 2 Department of Astronomy, Michigan State University, 567 Wilson Road, East Lansing, MI 48824-2320, USA 3 Physics Department and Institute of Theoretical and Computational Physics, University of Crete, 71003 Heraklion, Crete, Greece 4 Foundation of Research and Technology, 71003, Heraklion, Crete, Greece 5 Department of Astronomy & Astrophysics, 525 Davey Lab, The Pennsylvania State University, University Park, PA 16802, USA 6 Department of Physics & Astronomy, University of Leicester, Leicester, UK Received 2013 April 16; accepted 2013 June 18; published 2013 July 29

ABSTRACT We present a new method for the determination of the two-dimensional (2D) projected spatial distribution of globular clusters (GCs) in external galaxies. This method is based on the K-Nearest Neighbor density estimator of Dressler, complemented by Monte-Carlo simulations to establish the statistical significance of the results. We apply this method to NGC 4261, a “test galaxy” where significant 2D anisotropy in the GC distribution has been reported. We confirm that the 2D distribution of GC is not azimuthally isotropic. Moreover, we demonstrate that the 2D distribution departures from the average GC radial distribution results in highly significant spiral-like or broken shell features. Overall, the same perturbations are found in “red” and “blue” GCs, but with some differences. In particular, we observe a central feature, roughly aligned with the minor axis of NGC 4261, composed of red and most luminous GCs. Blue and fainter GCs are more frequent at large radial distances and follow the spiral-like features of the overall density structure. These results suggest a complex merging history for NGC 4261. Key words: galaxies: individual (NGC 4261) – globular clusters: general – methods: statistical Online-only material: color figures

1. INTRODUCTION of the peculiar observational traits of the GC populations, no agreement has yet been reached in the literature. A considerable body of work has been accumulated in the While the spectral and photometric properties of GC popula- past two decades on the properties of globular cluster (GC) tions, as well as their radial distributions, have been explored in populations in elliptical galaxies (see review by Brodie & depth, relatively little work has addressed their two-dimensional Strader 2006). The studies available in the literature that use (2D) distributions (e.g., NGC 4471 in Rhode & Zepf 2001, both Hubble Space Telescope (HST) and larger ground-based NGC 1399 in Dirsch et al. 2003; Bassino et al. 2006, NGC 3379, telescopes, have established the existence of red (metal rich) NGC 4406 and NGC 4594 in Rhode & Zepf 2004, NGC 4636 and blue (metal poor) subpopulations of GCs in most galaxies. in Dirsch et al. 2005, multiple galaxies in Hargis & Rhode 2012, The radial distributions of these populations differ, with the NGC 3585 and NGC 5812 in Lane et al. 2013). This type of red/metal rich GCs being more centrally concentrated than blue work requires us to extract the entire GC population with full GCs. These results suggest different formation scenarios for the or near-full coverage of the parent galaxy. Augmented by kine- two GC subpopulations, which may reflect the formation history matics of the GC systems and subsystems (Strader et al. 2011 of their parent galaxy. Brodie & Strader (2006), in particular, for M87; Blom et al. 2012 for NGC 4365), these studies are argue that blue GCs may have formed in low-mass dark matter bringing forth a picture of complex and diverse GC populations halos in the early universe, while the red GCs were built in elliptical galaxies, consistent with a history of gravitational in subsequent dissipative buildup of the parent galaxy. Other interactions and merging with neighboring galaxies. studies have investigated possible formation mechanisms for Here we study the 2D GC distribution of NGC 4261, for the families of low-metallicity and high-metallicity GCs using which deep observations of the innermost region of the galaxy, different approaches. These studies include semi-analytical obtained with HST, are reported by Bonfini et al. (2012, models that use galaxy assembly history from cosmological hereinafter B+12). NGC 4261 first came to attention because simulations and observed scaling relations to estimate the of the inhomogeneous spatial distribution of its low-mass X-ray amount and metallicity of cold gas available for star formation binaries (LMXBs), which suggested a possible minor merger (Muratov & Gnedin 2010), physical models of the collapse event (Zezas et al. 2003). Since a number of these LMXBs were leading to the formation of GCs based on high mass resolution associated with GCs (Giordano et al. 2005), the entire population simulations (Griffen et al. 2010), and models based on the of 718 GCs was extracted and studied with HST, revealing an observed galaxy mass–metallicity relation, the galaxy stellar azimuthal asymmetry in its distribution, which was attributed to mass function and theoretical merger rates (Tonini 2013). Other past minor merging or interaction (B+12). Ferrarese et al. (1996) authors have explained the formation of low-metallicity GCs also reported boxy isophotes which support the hypothesis of as starburst remnants of old dwarf galaxies that could have recent gravitational interaction experienced by NGC 4261. entered the halos of spiral galaxies (Elmegreen et al. 2012). Although the results of B+12 are convincing, our more While these works have explored multiple possible explanation advanced analysis methods provide a clearer picture of the

1 The Astrophysical Journal, 773:87 (12pp), 2013 August 20 D’Abrusco et al.

Figure 1. Left: distribution of the V − I color of the NGC 4261 GCs sample. The vertical thick line (V − I = 1.15) separates blue and red GCs, and the shaded gray area indicates the interval of color thresholds ([1.05, 1.3]) which determine density and residual maps for the two color classes qualitatively similar to the maps obtained with the (V − I = 1.15) color threshold (see discussion in Sections 2 and 3.1). Right: distribution of the I magnitude of the same sample. The vertical thick line separates high luminosity and low luminosity GCs and the shaded gray area indicates the interval of I magnitude value which determine density and residual maps similar to the maps obtained using the I = 23 threshold (Section 3.1). The shading lines region for I>24 has been not used to generate the density and residual maps for high-L and low-L GCs classes in Section 3.1. asymmetry and its statistical significance. Here we report the Table 1 re-analysis of the 2D distribution of this GC population aimed at Summary of the Samples of GCs Observed in NGC 4261 Used in the Paper quantifying the reality and shape of the spatial features in the 2D Ntot Nred Nblue NHighL NLowL GC distribution. We have used the K-Nearest Neighbor method (KNN; Dressler 1980) to identify the 2D features, supplemented NGC 4261 718 306 412 316 402(84) by Monte-Carlo simulations to test their statistical significance. Note. The number in parenthesis in the last column represent the GCs fainter In Section 2 the data used in this paper are described. The than I = 24 which are not used to produce the density and residual maps for method and the results of the its application to NGC 4261 are low-L GCs. discussed in Section 3. Our findings are discussed in Section 4 and summarized in Section 5. following analysis. Changing the color and luminosity bound- 2. DATA aries (shaded regions in the plots in Figure 1) does not affect our results. Details can be found in Section 3.2. We have used the B+12 catalog of GC positions and proper- We have excluded from our analysis the central circular region ties, which lists 718 GCs, within the D25 ellipse of NGC 4261 with r<0.42 where incompleteness in the detection of the GCs (de Vaucouleurs et al. 1991). From this sample, we have ex- is substantial (B+12). We have also excluded the undersampled tracted color and magnitude based subsamples. Figure 1 (left) regions outside the D25 isophote (de Vaucouleurs et al. 1991). shows the V − I histogram where blue (V − I<1.15) and red (V − I 1.15) GCs are separated following B+12. These au- 3. THE ANALYSIS thors found that V − I = 1.15 is the color corresponding to the 50% probability of the GC to belong to either the red or the blue For each of the GCs samples listed in Table 1,wehave subpopulations, assuming a two-Gaussian model of the color determined the 2D spatial distributions by applying the KNN distribution. However, as discussed in B+12, the GC color dis- density estimator (Dressler 1980). This density is based on the tribution does not show the clear bimodality typical of the GC local distribution of GCs, i.e., on the distances of the closest color distribution in other early-type galaxies (Brodie & Strader GCs. For each knot of a regular grid covering the region where 2006; Peng et al. 2006). Figure 1 (right) shows the histogram the density is to be determined, we measured the distance of of I magnitudes; we arbitrarily define high luminosity (high L) the Kth nearest neighbor GC (DK) from the position of the knot. GCs those with I<23 mag, and low luminosity (low L)GCs We used this approach in order to have density estimates even in those with I 23 mag. The I<23 value used to separate the low density regions of the GC distribution. The point-density is estimated as: low-L from high-L GCs was set to obtain equipopulated classes K of sources (see Table 1). The density and residual maps obtained DK = (1) for the luminosity classes defined using I<23 are described VD(DK ) in Section 3.1. Figure 2 shows the spatial distributions of GC where K is the index of the nearest neighbor used to calculate positions in the plane of the sky, where the azimuthal asym- the density; for example, for K = 5 only the five GCs nearest to metry in the 2D projected GC distribution is evident. Table 1 the grid knot are used. VD is the volume of the region within the lists the number of GCs in each of the main samples used in the distance DK of the Kth nearest neighbor from the point where

CONSTRAINTS ON THE GALACTIC POPULATION OF TeV PULSAR WIND NEBULAE USING FERMI LARGE AREA TELESCOPE OBSERVATIONS

F. Acero1, M. Ackermann2, M. Ajello3, A. Allafort4, L. Baldini5, J. Ballet6, G. Barbiellini7,8, D. Bastieri9,10, K. Bechtol4, R. Bellazzini11, R. D. Blandford4, E. D. Bloom4, E. Bonamente12,13, E. Bottacini4, T. J. Brandt1, J. Bregeon11, M. Brigida14,15, P. Bruel16, R. Buehler4,S.Buson9,10, G. A. Caliandro17, R. A. Cameron4, P. A. Caraveo18, C. Cecchi12,13, E. Charles4, R. C.G. Chaves6, A. Chekhtman19,58,J.Chiang4, G. Chiaro10,S.Ciprini20,21,R.Claus4, J. Cohen-Tanugi22, J. Conrad23,24,25,59,S.Cutini20,21, M. Dalton26,60, F. D’Ammando27, F. de Palma14,15, C. D. Dermer28, L. Di Venere4, E. do Couto e Silva4, P. S. Drell4, A. Drlica-Wagner4, L. Falletti22, C. Favuzzi14,15, S. J. Fegan16, E. C. Ferrara1, W. B. Focke4, A. Franckowiak4, Y. Fukazawa29, S. Funk4,P.Fusco14,15, F. Gargano15, D. Gasparrini20,21, N. Giglietto14,15, F. Giordano14,15, M. Giroletti27, T. Glanzman4, G. Godfrey4,T.Gregoire´ 30,31,I.A.Grenier6, M.-H. Grondin30,31, J. E. Grove28, S. Guiriec1, D. Hadasch17, Y. Hanabata29, A. K. Harding1, M. Hayashida4,32, K. Hayashi29, E. Hays1,J.Hewitt1,A.B.Hill4,33,61, D. Horan16,X.Hou26,R.E.Hughes34,Y.Inoue4, M. S. Jackson24,35, T. Jogler4,G.Johannesson´ 36,A.S.Johnson4, T. Kamae4,T.Kawano29, M. Kerr4,J.Knodlseder¨ 30,31,M.Kuss11, J. Lande4, S. Larsson23,24,37, L. Latronico38, M. Lemoine-Goumard26,60, F. Longo7,8, F. Loparco14,15, M. N. Lovellette28, P. Lubrano12,13, M. Marelli18, F. Massaro4, M. Mayer2, M. N. Mazziotta15,J.E.McEnery1,39, J. Mehault26,60, P. F. Michelson4, W. Mitthumsiri4,T.Mizuno40,C.Monte14,15, M. E. Monzani4, A. Morselli41, I. V. Moskalenko4, S. Murgia4, T. Nakamori42,R.Nemmen1,E.Nuss22,T.Ohsugi40, A. Okumura4,43, M. Orienti27, E. Orlando4,J.F.Ormes44, D. Paneque4,45, J. H. Panetta4, J. S. Perkins1,46,47,48, M. Pesce-Rollins11,F.Piron22, G. Pivato10, T. A. Porter4, S. Raino` 14,15, R. Rando9,10, M. Razzano11,49,A.Reimer4,50,O.Reimer4,50, T. Reposeur26,S.Ritz49, M. Roth51, R. Rousseau26,60, P. M. Saz Parkinson49, A. Schulz2, C. Sgro` 11,E.J.Siskind52,D.A.Smith26, G. Spandre11, P. Spinelli14,15,D.J.Suson53, H. Takahashi29, Y. Takeuchi42, J. G. Thayer4, J. B. Thayer4, D. J. Thompson1, L. Tibaldo4, O. Tibolla54, M. Tinivella11, D. F. Torres17,55, G. Tosti12,13,E.Troja1,62, Y. Uchiyama4, J. Vandenbroucke4, V. Vasileiou22, G. Vianello4,56, V. Vitale41,57, M. Werner50,B.L.Winer34,K.S.Wood28, and Z. Yang23,24 1 NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA 2 Deutsches Elektronen Synchrotron DESY, D-15738 Zeuthen, Germany 3 Space Sciences Laboratory, University of California, 7 Gauss Way, Berkeley, CA 94720-7450, USA 4 W. W. Hansen Experimental Physics Laboratory, Kavli Institute for Particle Astrophysics and Cosmology, Department of Physics and SLAC National Accelerator Laboratory, Stanford University, Stanford, CA 94305, USA; [email protected], [email protected] 5 Universita` di Pisa and Istituto Nazionale di Fisica Nucleare, Sezione di Pisa, I-56127 Pisa, Italy 6 Laboratoire AIM, CEA-IRFU/CNRS/Universite´ Paris Diderot, Service d’Astrophysique, CEA Saclay, F-91191 Gif sur Yvette, France 7 Istituto Nazionale di Fisica Nucleare, Sezione di Trieste, I-34127 Trieste, Italy 8 Dipartimento di Fisica, Universita` di Trieste, I-34127 Trieste, Italy 9 Istituto Nazionale di Fisica Nucleare, Sezione di Padova, I-35131 Padova, Italy 10 Dipartimento di Fisica e Astronomia “G. Galilei,” Universita` di Padova, I-35131 Padova, Italy 11 Istituto Nazionale di Fisica Nucleare, Sezione di Pisa, I-56127 Pisa, Italy 12 Istituto Nazionale di Fisica Nucleare, Sezione di Perugia, I-06123 Perugia, Italy 13 Dipartimento di Fisica, Universita` degli Studi di Perugia, I-06123 Perugia, Italy 14 Dipartimento di Fisica “M. Merlin” dell’Universita` e del Politecnico di Bari, I-70126 Bari, Italy 15 Istituto Nazionale di Fisica Nucleare, Sezione di Bari, 70126 Bari, Italy 16 Laboratoire Leprince-Ringuet, Ecole´ polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France 17 Institut de Ciencies` de l’Espai (IEEE-CSIC), Campus UAB, E-08193 Barcelona, Spain 18 INAF-Istituto di Astrofisica Spaziale e Fisica Cosmica, I-20133 Milano, Italy 19 Center for Earth Observing and Space Research, College of Science, George Mason University, Fairfax, VA 22030, USA 20 Agenzia Spaziale Italiana (ASI) Science Data Center, I-00044 Frascati (Roma), Italy 21 Istituto Nazionale di Astrofisica–Osservatorio Astronomico di Roma, I-00040 Monte Porzio Catone (Roma), Italy 22 Laboratoire Univers et Particules de Montpellier, Universite´ Montpellier 2, CNRS/IN2P3, F-34095 Montpellier, France 23 Department of Physics, Stockholm University, AlbaNova, SE-106 91 Stockholm, Sweden 24 The Oskar Klein Centre for Cosmoparticle Physics, AlbaNova, SE-106 91 Stockholm, Sweden 25 The Royal Swedish Academy of Sciences, Box 50005, SE-104 05 Stockholm, Sweden 26 Universite´ Bordeaux 1, CNRS/IN2p3, Centre d’Etudes´ Nucleaires´ de Bordeaux Gradignan, F-33175 Gradignan, France; [email protected], [email protected] 27 INAF Istituto di Radioastronomia, I-40129 Bologna, Italy 28 Space Science Division, Naval Research Laboratory, Washington, DC 20375-5352, USA 29 Department of Physical Sciences, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526, Japan 30 CNRS, IRAP, F-31028 Toulouse cedex 4, France 31 GAHEC, Universite´ de Toulouse, UPS-OMP, IRAP, F-31028 Toulouse, France 32 Department of Astronomy, Graduate School of Science, Kyoto University, Sakyo-ku, Kyoto 606-8502, Japan 33 School of Physics and Astronomy, University of Southampton, Highfield, Southampton SO17 1BJ, UK 34 Department of Physics, Center for Cosmology and Astro-Particle Physics, The Ohio State University, Columbus, OH 43210, USA 35 Department of Physics, Royal Institute of Technology (KTH), AlbaNova, SE-106 91 Stockholm, Sweden 36 Science Institute, University of Iceland, IS-107 Reykjavik, Iceland 37 Department of Astronomy, Stockholm University, SE-106 91 Stockholm, Sweden 38 Istituto Nazionale di Fisica Nucleare, Sezione di Torino, I-10125 Torino, Italy 39 Department of Physics and Department of Astronomy, University of Maryland, College Park, MD 20742, USA 40 Hiroshima Astrophysical Science Center, Hiroshima University, Higashi-Hiroshima, Hiroshima 739-8526, Japan 41 Istituto Nazionale di Fisica Nucleare, Sezione di Roma “Tor Vergata,” I-00133 Roma, Italy

1 The Astrophysical Journal, 773:77 (27pp), 2013 August 10 Acero et al.

42 Research Institute for Science and Engineering, Waseda University, 3-4-1 Okubo, Shinjuku, Tokyo 169-8555, Japan 43 Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya 464-8601, Japan 44 Department of Physics and Astronomy, University of Denver, Denver, CO 80208, USA 45 Max-Planck-Institut fur¨ Physik, D-80805 Munchen,¨ Germany 46 Department of Physics and Center for Space Sciences and Technology, University of Maryland Baltimore County, Baltimore, MD 21250, USA 47 Center for Research and Exploration in Space Science and Technology (CRESST) and NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA 48 Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA 49 Santa Cruz Institute for Particle Physics, Department of Physics and Department of Astronomy and Astrophysics, University of California at Santa Cruz, Santa Cruz, CA 95064, USA 50 Institut fur¨ Astro- und Teilchenphysik and Institut fur¨ Theoretische Physik, Leopold-Franzens-Universitat¨ Innsbruck, A-6020 Innsbruck, Austria 51 Department of Physics, University of Washington, Seattle, WA 98195-1560, USA 52 NYCB Real-Time Computing Inc., Lattingtown, NY 11560-1025, USA 53 Department of Chemistry and Physics, Purdue University Calumet, Hammond, IN 46323-2094, USA 54 Institut fur¨ Theoretische Physik and Astrophysik, Universitat¨ Wurzburg,¨ D-97074 Wurzburg,¨ Germany 55 Institucio´ Catalana de Recerca i Estudis Avan¸cats (ICREA), E-08010 Barcelona, Spain 56 Consorzio Interuniversitario per la Fisica Spaziale (CIFS), I-10133 Torino, Italy 57 Dipartimento di Fisica, Universita` di Roma “Tor Vergata,” I-00133 Roma, Italy Received 2013 April 15; accepted 2013 June 20; published 2013 July 29

ABSTRACT Pulsar wind nebulae (PWNe) have been established as the most populous class of TeV γ -ray emitters. Since launch, the Fermi Large Area Telescope (LAT) has identified five high-energy (100 MeV < E < 100 GeV) γ -ray sources as PWNe and detected a large number of PWN candidates, all powered by young and energetic pulsars. The wealth of multi-wavelength data available and the new results provided by Fermi-LAT give us an opportunity to find new PWNe and to explore the radiative processes taking place in known ones. The TeV γ -ray unidentified (UNID) sources are the best candidates for finding new PWNe. Using 45 months of Fermi-LAT data for energies above 10 GeV, an analysis was performed near the position of 58 TeV PWNe and UNIDs within 5◦ of the Galactic plane to establish new constraints on PWN properties and find new clues on the nature of UNIDs. Of the 58 sources, 30 were detected, and this work provides their γ -ray fluxes for energies above 10 GeV. The spectral energy distributions and upper limits, in the multi-wavelength context, also provide new information on the source nature and can help distinguish between emission scenarios, i.e., between classification as a pulsar candidate or as a PWN candidate. Six new GeV PWN candidates are described in detail and compared with existing models. A population study of GeV PWN candidates as a function of the pulsar/PWN system characteristics is presented. Key word: gamma rays: general Online-only material: color figures

1. INTRODUCTION (IC) scattering of highly relativistic leptons from the source on the ambient photon fields such as cosmic microwave background Since 2003, the extensive observations of the Galactic plane (CMB), stellar radiation, or infrared emission from dust (e.g., de by Cerenkov telescopes have detected more than 80 Galactic Jager et al. 2009; de Jager & Djannati-Ata¨ı 2009). This leptonic TeV sources (Hinton & Hofmann 2009). Pulsar wind nebu- approach could well explain several UNIDs (e.g., Tibolla 2011; lae (PWNe) are the dominant class with more than 30 firm Tibolla et al. 2012); moreover, one of its biggest advantages is identifications. A similar number of Galactic sources cannot be that it provides a natural explanation for the UNIDs that lack a associated with a counterpart at any other wavelength; they form lower energy (radio and X-ray) counterpart (e.g., de Jager et al. the unidentified (UNID) source class. The third largest class of 2009; H.E.S.S. Collaboration et al. 2012b), i.e., the so-called Galactic sources are the supernova remnants (SNRs). dark sources. In the hadronic scenario, hadrons accelerated by The Large Area Telescope (LAT) on board the Fermi Gamma- a source collide with the nuclei in the ambient medium (e.g., ray Space Telescope provides all-sky coverage of the γ -ray molecular cloud) and secondary neutral pions decay to γ -rays sky at energies from 20 MeV to more than 300 GeV. With (e.g., Gabici et al. 2009). 2 yr of observations, the Fermi-LAT Second Source Catalog The leptonic PWN scenario requires an energetic and young (2FGL; Nolan et al. 2012) reports the detection of 1873 sources, pulsar to be present. Pulsars are the largest class of Galactic 1298 being identified and 575 without clear identification. Four sources detected above 100 MeV with the LAT. These pulsars hundred of them lie within 5◦ of the Galactic plane. could make up part of the LAT UNID population. In the LAT Most of the LAT UNID sources are expected to be pulsars, energy range, pulsars are point-like sources and exhibit power- SNRs, binary systems, or PWNe. The γ -ray emission from law spectra with exponential cutoffs between 0.5 and 6 GeV these sources is expected to be either hadronic or leptonic. In the (Abdo et al. 2010e), while PWNe have hard power-law spectra leptonic scenario, γ -ray photons are created by inverse Compton without cutoffs in the GeV energy range and might be spatially resolved by the LAT. Middle-aged SNRs, interacting with 58 Resident at Naval Research Laboratory, Washington, DC 20375, USA. molecular clouds, detected by the LAT are generally bright and 59 Royal Swedish Academy of Sciences Research Fellow, funded by a grant exhibit a break at ∼2 GeV (Uchiyama 2011). Radio, X-ray, and from the K. A. Wallenberg Foundation. 60 γ -ray photons probe the non-thermal particle populations Funded by contract ERC-StG-259391 from the European Community. and therefore provide information to discriminate between sce- 61 Funded by a Marie Curie IOF, FP7/2007-2013, grant agreement No. 275861. narios in which the γ -ray emission is dominated by leptonic or 62 NASA Postdoctoral Program Fellow, USA. hadronic processes.

SHOCK EMERGENCE IN SUPERNOVAE: LIMITING CASES AND ACCURATE APPROXIMATIONS

Stephen Ro and Christopher D. Matzner Department of Astronomy & Astrophysics, University of Toronto, 50 St. George St., Toronto, ON M5S 3H4, Canada Received 2013 April 15; accepted 2013 June 25; published 2013 July 29

ABSTRACT We examine the dynamics of accelerating normal shocks in stratiﬁed planar atmospheres, providing accurate ﬁtting formulae for the scaling index relating shock velocity to the initial density and for the post-shock acceleration factor as functions of the polytropic and adiabatic indices which parameterize the problem. In the limit of a uniform initial atmosphere, there are analytical formulae for these quantities. In the opposite limit of a very steep density gradient, the solutions match the outcome of shock acceleration in exponential atmospheres. Key words: shock waves – supernovae: general

1. INTRODUCTION in turn, can be used to predict the amount and upper speed limit of the fastest ejecta and properties of the breakout flash (Matzner Shock emergence at the surface of an exploding star is & McKee 1999; Calzavara & Matzner 2004), transition to rel- an important moment in the life of a supernova. Shock and ativistic flow and aspects of the circumstellar interaction (Tan post-shock acceleration in the outer stellar envelope, and the et al. 2001), and many other breakout-related phenomena. With breakout of post-shock radiation from a thin layer beneath the advances in the theory of photon-mediated shocks and emission photosphere, can have a number of significant consequences. around the time of breakout (e.g., Katz et al. 2010, 2012; Nakar The escaping flash of radiation gives an energetic precursor & Sari 2010; Sapir et al. 2011), of the ultra-relativistic self- which can signal the supernova’s existence (Klein & Chevalier similar problem (Perna & Vietri 2002; Nakayama & Shigeyama 1978) and carries physical information about the explosion 2005; Pan & Sari 2006; Kikuchi & Shigeyama 2007), and of the (Matzner & McKee 1999; Calzavara & Matzner 2004; Nakar interaction of relativistic ejecta with a stellar wind (Nakamura & Sari 2012, 2010; Sapir et al. 2011;Katzetal.2010; Suzuki & Shigeyama 2006), among others, there are ample opportuni- & Shigeyama 2010; Piro et al. 2010); traveling outward, it can ties for these approximate global models to be improved and ionize a circumstellar nebula like the one surrounding SN 1987A extended. (Lundqvist & Fransson 1996) and produce an infrared light echo To advance this larger project, we focus here on the planar, as it encounters dust (Dwek & Arendt 2008). Shock emergence adiabatic, non-relativistic problem of an accelerating normal launches the fastest ejecta, the first to host the supernova shock. Our goal is to provide flexible yet highly accurate photosphere (Chevalier 1992) and the first to interact with approximations for the most important flow quantities, the shock circumstellar and interstellar matter, producing a synchrotron- acceleration index and the post-shock acceleration factor, as emitting shell (Fransson et al. 1996). If they meet a companion functions of the adiabatic and polytropic indices (γ and γp, star or dense circumstellar disk, then an additional X-ray signal respectively) which parameterize the problem. A secondary goal can be produced (Metzger 2010; Kasen 2010). is to demonstrate that although the flow quantities must typically In particularly compact and energetic explosions, the shock be found as eigenvalues of the dynamical problem, they adhere can become relativistic before emerging, and relativistic ejecta to well-understood limiting forms in several asymptotic cases. can create X-ray and γ -ray transients in their circumstellar The self-similar problem with a power-law atmosphere be- collisions (Matzner & McKee 1999; Tan et al. 2001), and low vacuum was posed by Gandel’Man & Frank-Kamenetskii may produce light elements through spallation (Fields et al. (1956) and solved in its Eulerian form by Sakurai (1960). We 2002; Nakamura & Shigeyama 2004) and, potentially, ultra- shall use Sakurai’s eigenvalue method to identify the shock ac- high-energy cosmic rays (Wang et al. 2007; Budnik et al. celeration index, but for the post-shock flow we employ the 2008). Lagrangian approach by Matzner & McKee (1999). This has Potentially observable shock breakouts accompany several the dual advantage that it continuously describes both the pre- other types of astrophysical events, including the type Ia breakout and post-breakout flow in a single function, and that it explosions (Piro et al. 2010) and accretion-induced collapses naturally connects each fluid element’s state at the shock front (Fryer et al. 1999; Tan et al. 2001) of white dwarfs, tidal with those in the final state. disruptions of stars, jet and cocoon emergence in long-duration gamma-ray bursts, and (albeit in a less energy-conserving 2. PROBLEM, METHOD, AND SOLUTIONS manner) superbubbles in galactic disks. Underlying all these phenomena are the hydrodynamics of Our problem involves one-dimensional flow with an altitude shock acceleration in the outer layers of a star, and anchor- x relative to the stellar surface. The initial density distribution of n ing these dynamics is the asymptotic problem of flow be- cold matter is ρ0(x) ∝ (−x) for x<0 and ρ0(0) = 0forx>0. hind a normal, adiabatic shock accelerating through a planar Here, n is the polytrope parameter, which is related to γp by the medium which varies as a power law with depth. As Matzner & hydrostatic relation with constant gravity g∗:ifP = 0atx = 0, γp McKee (1999) first demonstrated, this asymptotic planar solu- then P (x) = g∗(−x)ρ(x)/(n +1)∝ ρ(x) with γp = 1+1/n. tion can be combined with the dynamics of a spherical, self- A strong adiabatic shock wave accelerates down this density similar blastwave into an accurate approximate model for shock gradient, reaching x = 0att = 0 with infinite velocity; propagation and post-shock flow in a spherical explosion. This, neglecting radiative effects, this is the point of breakout.

1 The Astrophysical Journal, 773:79 (5pp), 2013 August 10 Ro & Matzner For t>0 the shock disappears and matter expands into the equation. In general, this is only a good approximation, because region of positive x. In the limit that all additional physical the shock moves more slowly than these forward characteristics effects—curvature, gravity, and temperature of the star, finite and because conditions vary from one characteristic to another. depth, non-simultaneity of breakout, relativity, shock thickness, In the limit n → 0, however, there is no variation among etc.—are negligible, the flow is self-similar. The shock velocity characteristics, and Whitham’s approximation becomes exact. −λ −β accelerates according to vs =˙xs (t) ∝ (−x) ∝ ρ , where The isothermal limit γ → 1 is characterized by β = λ/n. The fluid motion is a universal function of self- | | similar variables like x/xs ( t )orx(m, t)/x0(m), in which each β → 0 · (γ → 1). fluid element (labeled by its mass coordinate m) accelerates from a post-shock velocity toward its terminal velocity, which In this limit, a strong shock is infinitely compressive and is a unique multiple vf (m)/vs(m) of the shock velocity which governed by the conservation of momentum. The shock velocity crossed that element. Our task is to find β and vf /vs as functions is constant because the atmosphere above the shock front has of γ and γp. negligible mass relative to the shell of post-shock material. We note that Whitham’s approximation is exact in this limit as well, 2.1. Shock Acceleration Parameter and Its Limits because the shock no longer outruns forward characteristics. →∞ To find the shock acceleration index λ, or equivalently the In the limit n , the stellar structure is isothermal → velocity-density index β, we follow Sakurai (1960). Sakurai (γp 1) and transitions from power-law to exponential in its writes the conservation equations for mass, energy, and entropy depth dependence. It is reassuring, therefore, that Equation (1) in Eulerian form, introduces the self-similar ansatz, and arrives gives at a single, first-order differential equation for the spatial −1 structure of the post-shock flow prior to breakout. This equation 4.321γ 0.469 must pass smoothly from the conditions immediately behind β → . n →∞ 1 78 + − ( ) the shock front, through a critical point at the sonic point γ 1 of the flow; this is only possible for a unique value of λ in this limit. As Hayes (1968) notes, this limit coincides with the or β, which we identify by a shooting method. We present case of an exponential atmosphere (γ = 1); we reproduce his =− p the solution space β(n, γ ) spanning log10 n 6,...,6 and solutions and those of Raizer (1964). Equation (1) demonstrates − =− log10(γ 1) 5,...,6inTable1. For the entire parameter that the shock acceleration index is indistinguishable from the space, the simple functional fit exponential case for all n 102—or in practical terms, any −1 time that the distance to the surface is very far when measured Bγ C relative to the density scale height. β = A + with (1) →∞ γ − 1 In the γ limit of incompressible flow, β takes the definite form [A(n)+B(n)C(n)]−1. We know of no physical 0.22 A = 2 − , explanation for this result. 0.59n−9/8 +1 We end this section by noting that for the specific case 2.31 γ = 3/2 and n = 5, we find β = 1/5 and λ = 1 (at least B = 2+ , and 8 −1 to within a part in 10 , while the fit of Equation (1)gives 1.72n +1 = 1 0.0312 λ 1.01). This is unlikely to be a coincidence, although we C = − have not identified any simplification in the dynamical equations −1 2 1.1n +1 for this case. is quite accurate: the root-mean-square (rms) error relative to the values in Table 1 is 1.4%, and the error, which is concentrated at 3. POST-SHOCK FLOW AND ASYMPTOTIC high n and low γ , is at most 4.0%. High accuracy is necessary, FREE EXPANSION because β is the exponent of a number which becomes large around breakout. For instance, the energy in relativistic ejecta After each parcel of gas has been swept into motion by the ∝ 2 γp/(2β) scale as Erel [Ein/(Mejc )] , where Ein is the explosion shock, it continues to accelerate until its internal energy is spent energy and M is the total ejected mass; in the model for SN ej and it has reached the terminal velocity vf (m). To describe 1998bw discussed by Tan et al. (2001), an error of β leads to this, we employ the Lagrangian method of Matzner & McKee an error in E which is nine times greater. Higher accuracy rel (1999). This naturally provides quantities like vf (m)/vs (m), can be obtained by interpolating our table, and the differential and continues smoothly through the point of breakout; however, equations yield solutions to numerical accuracy. Several limiting it does not yield eigenvalues like β as readily as Sakurai’s forms of our fit to β(n, γ ) are readily apparent. method. In order to correct a couple typos in Matzner & → In the limit n 0 of an effectively uniform stellar envelope, McKee’s Appendix (which do not affect their results), we β and its fit in Equation (1) reproduce the approximate ex- write out the equation. The Lagrangian self-similar time and pression derived by Whitham (1958 and 1974, simplifying and space coordinates are η = t/t0(m) and S = x/x0(m), where improving upon results by Chisnell 1955, 1957, and Chester t0(m) is the time at which the shock crosses x0(m). For 1960): a given fluid element both η and S decline from unity to −1 2γ 1/2 −∞ as a fluid element accelerates outward; shock breakout β → 2+ · (n → 0). = γ − 1 (neglecting radiative effects) is at η 0, and the element exits the boundaries of the progenitor somewhat later, when Whitham arrived at this form by reasoning that quantities just S = 0. As Matzner & McKee discuss, the pressure and behind the shock front should evolve similarly to those found density distributions and the resulting acceleration can be along a forward-traveling sound wave, for which there is an exact computed from these variables, and the equating resulting fluid

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Skyrme effective pseudopotential up to the next-to-next-to-leading order

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Please note that terms and conditions apply. IOP PUBLISHING JOURNAL OF PHYSICS G: NUCLEAR AND PARTICLE PHYSICS J. Phys. G: Nucl. Part. Phys. 40 (2013) 095104 (8pp) doi:10.1088/0954-3899/40/9/095104

Skyrme effective pseudopotential up to the next-to-next-to-leading order

D Davesne1,APastore2 and J Navarro3

1 Universite´ de Lyon, F-69622 Lyon, France, Universite´ Lyon 1, Villeurbanne; CNRS/IN2P3, UMR5822, IPNL 2 Institut d’Astronomie et d’Astrophysique, CP 226, Universite´ Libre de Bruxelles, B-1050 Bruxelles, Belgium 3 IFIC (CSIC-Universidad de Valencia), Apartado Postal 22085, E-46071-Valencia, Spain

E-mail: [email protected]

Received 5 July 2013 Published 29 July 2013 Online at stacks.iop.org/JPhysG/40/095104

Abstract The explicit form of the next-to-next-to-leading order (N2LO)ofthe Skyrme effective pseudopotential compatible with all required symmetries and especially with gauge invariance is presented in a Cartesian basis. It is shown in particular that for such a pseudopotential there is no spin–orbit contribution and that the D-wave term suggested in the original Skyrme formulation does not satisfy the invariance properties. The six new N2LO terms contribute to both the equation of state and the Landau parameters. These contributions to symmetric nuclear matter are given explicitly and discussed.

Communicated by Jacek Dobaczewski

1. Introduction

Since the implementation of Vautherin and Brink [1], the Skyrme interaction [2, 3] has become a popular tool for the description of nuclear properties based on a self-consistent mean-field approach. In its standard form it consists of zero-range central, spin–orbit and density-dependent terms, involving up to 11 parameters, which are fitted to properties of infinite nuclear matter and some selected nuclei. Systematic studies of binding energies and one-body properties of nuclei can thus be successfully performed in a very wide region of the nuclear chart [4]. However, there are some nuclear properties which cannot be correctly reproduced with the standard Skyrme terms, especially as one moves away from the valley of β-stability [5]. Intense work is currently being devoted to improving the existing parametrizations, in particular considering additional terms to the standard form [6]. In the original proposal, an effective pseudopotential was constructed as an expansion of the nuclear interaction in relative momenta, thus simulating finite-range effects in a zero-range interaction. It actually contained more terms than the current standard form, which is limited to the second order in momenta plus a density-dependent term which comes from a zero-range

0954-3899/13/095104+08$33.00 © 2013 IOP Publishing Ltd Printed in the UK & the USA 1 The Astrophysical Journal, 773:96 (27pp), 2013 August 20 doi:10.1088/0004-637X/773/2/96 C 2013. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

PROBING OUR HELIOSPHERIC HISTORY. I. HIGH-RESOLUTION OBSERVATIONS OF Na i AND Ca ii ALONG THE SOLAR HISTORICAL TRAJECTORY

Katherine Wyman1,2 and Seth Redﬁeld2 1 Harvard-Smithsonian Center for Astrophysics, 60 Garden Street, Cambridge, MA 02138, USA; [email protected] 2 Astronomy Department, Van Vleck Observatory, Wesleyan University, Middletown, CT 06459, USA; sredﬁ[email protected] Received 2012 November 26; accepted 2013 June 3; published 2013 July 29

ABSTRACT Over the course of its motion through the Galaxy, our solar system has encountered many interstellar environments of varying characteristics. Interstellar medium (ISM) density variations spanning seven orders of magnitude are commonly seen throughout the general Galactic environment, and a sufficiently dense cloud within this range has the potential to dramatically alter the structure of the heliosphere. We present observations of the ISM environments the Sun has most recently encountered based on high-resolution optical spectra toward nearby stars in the direction of the historical solar trajectory. The data were obtained with the highest-resolution spectrographs available, including the Tull Spectrograph on the Harlan J. Smith Telescope at McDonald Observatory and the Ultra-High-Resolution Facility on the Anglo-Australian Telescope at the Anglo-Australian Observatory. Observations were made of interstellar Na i and Ca ii doublet absorption toward 43 bright stars within ∼500 pc. No absorption is seen out to a distance of 120 pc (consistent with the outer boundary of the Local Bubble), but a complex collection of absorbers is seen in stars beyond 120 pc. While common absorbers are consistently seen in most sight lines, significant spatial variation is also detected, even between closely spaced sight lines. This pervasive evidence of small-scale structure not only speaks to the complexity of the morphology or physical properties of the gas in the ISM, but also emphasizes that dramatic structural changes to the heliosphere are common and it is important to understand the implications of such changes, such as the modulation in the cosmic ray flux, on planets. Key words: ISM: clouds – ISM: structure – line: profiles – local interstellar matter – Sun: heliosphere – techniques: spectroscopic Online-only material: color figures

1. INTRODUCTION the most dramatic driver of the structure of the heliosphere originates outside its boundaries (i.e., the properties of the As stars hurtle along their orbits around the Galaxy, they pass surrounding ISM). through a diverse array of interstellar medium (ISM) environ- Currently the Sun is moving through a large (R ∼ 100 pc), − ments, ranging from dense molecular clouds (e.g., molecular hot (T ∼ 106 K), and low density (n ∼ 0.005 cm 3) cavity knots in Orion, n 105 cm−3; O’dell 2001) to rarefied regions, known as the Local Bubble. Found within the Local Bubble are a − heated and largely vacated by strong supernovae or stellar winds collection of smaller (1–10 pc), higher density (n ∼ 0.1cm 3), (as seems to be the case with the Local Bubble, n 10−2 cm−3; cooler (T ∼ 7000 K), partially ionized clouds (Redfield & Welsh & Shelton 2009). Such a dramatic variation in the char- Linsky 2008). Given current solar velocity measurements (e.g., acteristics of the local interstellar environment of stars can have Dehnen & Binney 1998), a cloud 1 pc in diameter would pass a profound impact on the conditions in the volume occupied by over the solar system in just 70,000 yr. Currently, the Sun resides planets. A balance of the pressures, between an outward mov- very near to the inside edge of one of these clouds referred to as ing stellar wind and the inward confinement of the surrounding the Local Interstellar Cloud (LIC). Given the relative velocity interstellar material, defines an interface, the boundary between vectors of the Sun and the LIC, the Sun will be transitioning out the star and the rest of the universe. In the case of our Sun, this of the LIC sometime in the next few thousand years (Redfield interface, or heliopause, is located at approximately 300 AU &Linsky2000). In fact, the farthest Oort Cloud objects might (e.g., Muller¨ et al. 2006), a direct consequence of the properties have already breached this boundary and moved into hot Local of the local interstellar medium (LISM) and the solar wind (i.e., Bubble material. The Local Bubble is also home to clouds of their density and velocity). much higher densities and cooler temperatures. The Local Leo Changes in the extent of the heliosphere occur over a Cold Cloud, a thin cloud first identified in H i 21 cm emission, multitude of time scales. In the short term, the solar wind has recently been identified to be located between 11.3 and strength fluctuates over its 22-yr magnetic cycle while solar 24.3 pc away (Meyer et al. 2006; Peek et al. 2011). This cloud flares and coronal mass ejections can drive variability on appears to have a temperature of ∼20 K and a H i density of − even shorter time scales. Coronal activity has the potential to ∼3000 cm 3 (Meyer et al. 2012). In addition, several of the asymmetrically modify the morphology of the heliosphere. An MBM clouds (Magnani et al. 1985), originally identified by event like this occurred in 2003 and was recorded by instruments 2.6 mm CO emission, may be at or within the Local Bubble aboard Voyager 1 (Decker et al. 2005). On longer time scales, boundary. In particular, in roughly the same direction as the a gradual change in the solar wind strength will occur over the survey presented here, high-resolution observations of nearby 10-billion-yr life span of the Sun. Observations of other solar stars were used to pinpoint the distance of MBM 20 (Penprase type stars suggest that the historical solar wind may have been 1993; Hearty et al. 2000). A comprehensive review of the LISM ∼50× stronger than it is today (Wood et al. 2005). However, is provided by Frisch et al. (2011).

1 The Astrophysical Journal, 773:96 (27pp), 2013 August 20 Wyman & Redfield Given the pliable nature of the heliosphere and the configura- discussion to Paper II, including the likelihood that the cur- tion of our local galactic environment, several questions arise: rent interstellar column in the direction of the historical solar How will the heliosphere respond when the Sun exits the LIC? motion is identical with that of the actual interstellar environ- What would happen if the Sun were to traverse a dense inter- ments traversed by the Sun over that time. In this paper, we stellar cloud, such as the Local Leo Cold Cloud? The current focus primarily on presenting the rich data set and providing an paper develops the foundation for such an investigation by ask- analysis of the ISM in this unique direction. ing: What were the interstellar environments of the Sun in its most recent past? 2. OBSERVATIONS The implications of a dynamic heliosphere go well beyond structural changes and could be important for planets and Stars were observed to measure ISM absorption in the their atmospheres. The relationship between our interstellar direction of the Sun’s path through the Galaxy, out to a distance environment and the Earth has interested researchers for almost of ∼500 pc. The direction and velocity of the solar peculiar a century. Shapley (1921) first formally suggested a link between velocity were derived from Dehnen & Binney (1998), who the passage of the solar system through dense ISM clouds measure the solar peculiar velocity from the Hipparcos data and quasi-periodic climate catastrophes and mass extinctions set (see Schonrich¨ et al. 2010 for an alternative analysis). ◦ ◦ of organisms on Earth. As a testament to its broad appeal, The adopted values are l0 = 207.70 and b0 =−32.41, theoretical speculations on this relationship have continued to in the constellation Eridanus. The magnitude of the solar this day (Hoyle & Lyttleton 1939; Fahr 1968; Begelman & Rees peculiar velocity is estimated to be 13.38 ± 0.42 km s−1, which 1976; Talbot et al. 1976; Frisch & York 1986; Thaddeus 1986; corresponds to a time of ∼36.7 million yr to traverse 500 pc. This Zank & Frisch 1999; Shaviv 2003). A recent compendium of represents less than 1% of the total path length the Sun traverses articles related to the solar journey through the Galaxy is given in a single orbit around the Galaxy. This small percentage allows in Frisch (2006). Mechanisms invoked to link the ISM and us to approximate the Sun’s recent historical trajectory as a planetary atmospheres include: (1) fluctuations of cosmic ray straight line toward l0, b0. flux modulated through the heliosphere, given that the solar Stars for this study were chosen based on a close proximity ◦ magnetic field extends only as far as the termination shock (θ<10 )tol0, b0, and properties that optimize ISM absorption (e.g., Giacalone & Jokipii 1999; Usoskin et al. 2005; Smith analysis (e.g., bright [V<8], rapidly rotating, early type & Scalo 2009); and (2) dust accretion directly onto planetary [B, A, F] stars). We use the revised Hipparcos parallaxes atmospheres (e.g., Pavlov et al. 2005). These processes have presented by van Leeuwen (2007) to calculate distances to our potential consequences for climate, ozone chemistry, and DNA target stars, and restricted our sample to targets for which the mutation rates for surface organisms (see review by Scherer relative error in distance is 0.5, whereas the median relative et al. 2006). error is 0.1. This restriction led us to remove three targets This work builds on a legacy of high-resolution ISM observa- (HD 30020, HD 29173, and HD 30076). tions in the optical. The strength of the Na i and Ca ii transitions We observed two targets (HD 26574, 37.31 pc; HD 26994, has made them invaluable tracers of the ISM, even when they 289 pc) that were more than 9 deg from the direction of the are trace ions in interstellar gas (as is the case in the LISM, historical solar trajectory. HD 26574, the closer target showed where for example, Ca iii is the dominant form of calcium and no interstellar absorption, while HD 26994 had absorption that Na ii is the dominant form of sodium; Slavin & Frisch 2008). was dramatically different from what was observed for the rest Surveys have been made along sight lines toward most of the of our sample. For this reason, we decided to constrain our bright stars across the entire sky (Vallerga et al. 1993; Welty sample cone to <9◦ from the solar historical trajectory. A third et al. 1994, 1996), of star-forming regions (e.g., Orion; Frisch target, HD 32964 (66 Eri), is a well-known short period binary et al. 1990; Welsh et al. 2005), and of wide binaries used to (P = 5.5 days; Makaganiuk et al. 2011) in which both stars search for small-scale structure in the ISM along almost iden- displayed stellar absorption features at velocities at which we tical sight lines (e.g., Meyer & Blades 1996). However, as far would expect to see LISM absorption. Because this target is as we know, this is the densest high-resolution absorption line only 94.7 pc away, and given the results from our entire sample, survey of any particular direction through the ISM. Not only is we would not expect to see any interstellar absorption since the this particular sight line of interest to understand the history of closest target with detected absorption is 125.9 pc and 112.9 pc solar interactions with the surrounding interstellar gas, but also for Ca ii and Na i, respectively. Therefore, we removed these this survey provides a detailed and densely sampled look at the targets from our sample. The remaining 43 targets are listed in physical properties of the galactic ISM. Table 1 and their distribution on the sky is illustrated in Figure 1. A companion paper (K. Wyman & S. Redfield 2013, in prepa- Data were obtained throughout the time period of 2003 Oc- ration, hereafter Paper II) explores the potential consequences tober through 2005 October. Absorption lines of neutral sodium an encounter with the ISM environments presented in this paper (Na i D1 and D2; 5895.9242 and 5889.9510 Å) and singly ion- could have on the solar system. Models of the heliosphere have ized calcium (Ca ii H and K; 3968.4673 and 3933.6614 Å) were been made to explore a wide range of different interstellar en- chosen as the ISM component tracers since they both have strong vironments by Muller¨ et al. (2006). In particular, models were doublet resonance lines in the optical that are sensitive to nearby made using characteristics of the clouds in the LISM, as well as interstellar clouds. Observations were made with the highest- the hot Local Bubble and cold dense molecular clouds. These resolution spectrographs available in each hemisphere. In the models provide relations between the interstellar properties and northern hemisphere, observations were taken with the Harlan heliospheric properties. Here, we present observations of the J. Smith Telescope, a 2.7 m telescope located at McDonald interstellar clouds traversed by the Sun and plan to use these Observatory in west Texas. Observations were made at three relationships to make estimates of the historical heliospheric separate resolution (R ≡ λ/Δλ) settings: TS12: Tull Spec- response (and thereby the cosmic ray flux history on the Earth’s trograph (R ∼ 400,000; Tull 1972), TS21: Cross-Dispersed atmosphere) in our companion paper. We defer additional Echelle Spectrograph (2d coude)´ focus 1 (R ∼ 240,000;

THE ROCHE LIMIT FOR CLOSE-ORBITING PLANETS: MINIMUM DENSITY, COMPOSITION CONSTRAINTS, AND APPLICATION TO THE 4.2 hr PLANET KOI 1843.03

Saul Rappaport1, Roberto Sanchis-Ojeda1, Leslie A. Rogers2,4, Alan Levine3, and Joshua N. Winn1 1 Department of Physics, and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA; [email protected], [email protected], [email protected] 2 Department of Astronomy and Department of Planetary Science, California Institute of Technology, MC 249-17, Pasadena, CA 91125, USA; [email protected] 3 37-575 MIT Kavli Institute for Astrophysics and Space Research, 70 Vassar Street, Cambridge, MA 02139, USA; [email protected] Received 2013 June 3; accepted 2013 July 12; published 2013 July 29

ABSTRACT The requirement that a planet must orbit outside of its Roche limit gives a lower limit on the planet’s mean density. The minimum density depends almost entirely on the orbital period and is immune to systematic errors in the stellar properties. We consider the implications of this density constraint for the newly identified class of small planets with periods shorter than half a day. When the planet’s radius is accurately known, this lower limit to the density can be used to restrict the possible combinations of iron and rock within the planet. Applied to KOI 1843.03, a −3 0.6 R⊕ planet with the shortest known orbital period of 4.245 hr, the planet’s mean density must be 7gcm .By modeling the planetary interior subject to this constraint, we find that the composition of the planet must be mostly iron, with at most a modest fraction of silicates (30% by mass). Key words: instabilities – planetary systems – planets and satellites: detection – planets and satellites: individual (KOI 1843) Online-only material: color figures

1. INTRODUCTION allowed orbital periods for planets of various iron–silicate compositions. There is a growing list of exoplanets with very short orbital periods, including about 20 candidates with periods shorter than 2. THE ROCHE LIMITING ORBITAL PERIOD half a day (Batalha et al. 2011; Muirhead et al. 2012; Rappaport et al. 2012; Ofir & Dreizler 2013; Huang et al. 2013; Sanchis- The Roche limiting distance for a body comprised of an Ojeda et al. 2013). Essentially all of these shortest-period planets incompressible fluid with negligible bulk tensile strength in a have radii smaller than 2 R⊕. There are a number of reasons circular orbit about its parent star is why larger planets are not likely to survive in such short-period 1/3 orbits. Among the perils of being a short-period gas giant are ρs tidally induced orbital decay (Rasio & Ford 1996), a possible amin 2.44R , (1) ρ tidal-inflation instability (Gu et al. 2003a), Roche lobe overflow p (Gu et al. 2003b), and evaporation (Murray-Clay et al. 2009). where ρ and ρ are the mean densities of the parent star and An Earth-mass rocky planet would be less susceptible to these s p of the planet, respectively, and R is the radius of the parent effects, and in particular the solid portion of the planet could star (Roche 1849). This is the familiar form of the Roche limit. survive evaporation nearly indefinitely (Perez-Becker & Chiang In cases where the orbital period P of the planet is measured 2013). directly it is more useful to rewrite the equation using Kepler’s However, even small planets must orbit outside of the Roche third law, (2π/P)2 = GM/a3. The stellar mass and radius limit, the distance within which the tidal force from the star cancel out, giving would disrupt the planet’s hydrostatic equilibrium and cause it to rapidly disintegrate. In this work, we show that the periods of − some of the newly identified planet candidates are so short that 3π(2.44)3 ρ 1/2 P 12.6hr p , (2) the Roche limit leads to astrophysically meaningful constraints min −3 Gρp 1gcm on the planet’s mean density. If, in addition, the planet’s radius can be accurately determined from the measured transit depth where P is the minimum orbital period that can be attained and the estimated stellar radius, the composition of the planet min by the planet before being tidally disrupted. Note that Pmin can be constrained as well. is essentially independent of the properties of the parent star We apply this technique to the transiting planet candidate (except if it is rapidly rotating and substantially oblate), and it with the shortest known period, which was reported by Ofir depends only on the density of the planet. Turned around, this & Dreizler (2013) and independently identified in our search states that for a given orbital period, there is a well defined lower for short-period planets in the Kepler database (R. Sanchis- limit on the planet’s mean density. Ojeda et al., in preparation). We find that the planet has a mean −3 For a planet comprised of a highly compressible fluid, the density 7gcm , a mass of about 1/2 M⊕, and a composi- mean radius of the Roche lobe of the body is tion likely dominated by iron. We also compute the shortest m 1/3 r 0.49 a, (3) 4 Hubble Fellow. L M

1 The Astrophysical Journal Letters, 773:L15 (5pp), 2013 August 10 Rappaport et al.

1 Planet Mean Density 10

1.5 ) −3 ) ) ⊕ ⊕ (R p 0 (R 1 p > (g cm 10 R R ρ < Planet Radius 0.5

−1 10 −2 −1 0 1 0 10 10 10 10 2 3 4 5 6 7 M (M ) p ⊕ Minimum Allowed Period (hours) Figure 1. Left: radius and mean density as a function of mass, for planets with an iron core and silicate mantle. The colors indicate the composition, ranging from pure iron (dark blue) to pure silicates (light blue), with steps of 0.05 in the silicate fraction. Right: planet radius vs. the Roche limiting minimum orbital period, computed from the radius–mass relations given in the left panel and Equation (5). The color coding is the same as for the left panel. The properties of KOI 1843.03 are shown with a black vertical line, with an extent indicating the uncertainty in the measured radius. The disjoint curve for a pure silicate composition results from the abrupt change in central density when the Fe core disappears, in conjunction with the prescription used in Equation (5). More generally, for small Fe cores, a simple density contrast ratio does not adequately represent the planet’s interior structure nor does Equation (5) accurately yield Pmin. (A color version of this figure is available in the online journal.) where m and M are the planet and star masses, respectively, 3. LIMITING PERIOD AND COMPOSITION and m M (Eggleton 1983). This can be translated into an expression analogous to Equation (2): Given a lower limit on mean density, we can restrict the possibilities for the planet’s composition with recourse to the- −1/2 oretical models for the planet’s interior. Here we use planet 3π ρp Pmin 9.6hr . (4) mass–radius relations computed with an approach similar to (0.49)3Gρ 1gcm−3 p that of Seager et al. (2007) and Sotin et al. (2007). We con- Planets composed of iron and silicates are neither of uniform sider planets comprised of an iron core and a silicate man- density nor highly compressible. For this regime, we have tle. The iron core is described using the equation of state for used the results of Lai et al. (1993), calculated for a range -phase Fe from Seager et al. (2007). For the mantle, we as- of polytropes, to derive a simple, approximate interpolating sume a distilled mineralogical and elemental make-up following formula between Equations (2) and (4): Sotin et al. (2007); trace elements are neglected, and only the dominant mineral phases combining the four most abundant el- −1/2 −0.16 ρp ρ0p ements in Earth’s mantle (Si, Mg, Fe, O) are considered. At low P 12.6hr , (5) min −3 1gcm ρp pressures, the mantle consists of olivine ([Mg, Fe]2SiO4) and pyroxene ([Mg, Fe]2Si2O6), while at pressures above ∼25 GPa, where ρ0p is the central density of the planet. For the planets the mantle is a mixture of perovskite ([Mg, Fe]SiO3) and mag- of interest, the ratio of central density to mean density (ρ0p/ρp) nesiowustite¨ ([Mg, Fe]O). We adopt a solar Si/Mg molar abun- 5 ranges between about 1 and 2.5. Equation (5) is valid only dance ratio (Si/Mg = 1.131) and Earth-like magnesium number for ρp0/ρp 6, at which point Equation (4) is recovered. We (Mg/[Mg + Fe]silicates = 0.9). believe this interpolating formula to be accurate for the bodies The left panel of Figure 1 shows the calculated mean density considered here, and use it for the remainder of this work. A and radius as a function of mass for planets of varying iron priority for future work is to derive a more exact expression fractions. The right panel shows the minimum orbital period for the case of a planet of arbitrary composition and central according to Equation (5). These figures demonstrate that for concentration. planets with orbital periods shorter than about 6 hr, the Roche We see from the above expressions that for planets with peri- limit places an astrophysically meaningful constraint on the ods shorter than about 12 hr, the Roche limit leads immediately interior composition. to constraints on any gaseous component, and for periods shorter than about 6 hr the Roche limit begins to be highly constrain- 4. APPLICATION TO KOI 1843.03 ing even for terrestrial bodies. Planets with periods as short as − 4 hr must have densities 7–9 g cm 3. For realistic equations As a concrete example we consider KOI 1843 (KIC 5080636), of state applied to Earth-sized planets, we will show that these a Kepler star with a 4.2 hr transit candidate identified by Ofir & densities imply a largely metallic composition, with a possible Dreizler (2013). Two longer-period transit signals had already modest layer of silicates. been identified for this target star by the Kepler team. We refer to the 4.2 hr signal as KOI 1843.03. This object provides a good 5 These density contrast ratios are much closer to unity than they are for stars. Lower main-sequence stars can be characterized as n = 3/2 polytropes for case study because it has the shortest known orbital period of which ρ0/ρ ≈ 6; upper main-sequence stars are close to n = 3 polytropes any transit candidate, and is very likely to be a bona fide planet ≈ for which ρ0/ρ 54. as discussed below.

KECK-I MOSFIRE SPECTROSCOPY OF THE z ∼ 12 CANDIDATE GALAXY UDFj-39546284∗

P. Capak1, A. Faisst2,J.D.Vieira3, S. Tacchella2, M. Carollo2, and N. Z. Scoville3 1 Spitzer Science Center, 314-6 Caltech, Pasadena, CA 91125, USA 2 Institute for Astronomy, Swiss Federal Institute of Technology (ETH Zurich), CH-8093 Zurich, Switzerland 3 California Institute of Technology, 314-6 Caltech, Pasadena, CA 91125, USA Received 2013 June 10; accepted 2013 July 12; published 2013 July 29

ABSTRACT We report the results of deep (4.6 hr) H-band spectroscopy of the well studied z ∼ 12 H-band dropout galaxy candidate UDFj-39546284 with MOSFIRE on Keck-I. These data reach a sensitivity of 5–10 × 10−19 erg s−1 cm−2 per 4.4 Å resolution element between sky lines. Previous papers have argued that this source could either be a large equivalent width line emitting galaxy at 2

1. INTRODUCTION and to quantify how star-forming galaxies contribute to the re- ionization of the Universe (Robertson et al. 2010; Bouwens et al. Deep photometric surveys with the Hubble Space Telescope 2011a; Oesch et al. 2012). (HST) are revolutionizing our knowledge of the star-forming However, spectroscopic redshifts are fundamental to ascertain galaxy population at the epoch of hydrogen re-ionization, only the true redshifts of candidate high-z LBGs. At z>4, the a few hundred million years after the big bang. Hundreds of application of the dropout technique has lead to notably diverse candidate z ∼ 7–8 galaxies have now been discovered using results (Iwata et al. 2003; Ouchi et al. 2004; Bouwens et al. the Lyman break galaxy (LBG or “dropout”) technique using 2007; van der Burg et al. 2010) with only limited samples newly available HST Wide-Field Camera 3 Infrared (WFC3- being spectroscopically confirmed (Vanzella et al. 2009;Stark IR) data (Bouwens et al. 2010; Oesch et al. 2010; McLure et al. 2010; Mallery et al. 2012), and the vast majority of et al. 2013; Ellis et al. 2013). Recently, this technique has been spectroscopy failing to detect anything significant. For example, pushed to z ∼ 12 by using hundreds of HST orbits to probe Stark et al. (2010) and Vanzella et al. (2009), two of the largest to unprecedented depths in the Hubble Ultra-Deep Field 2012 spectroscopic samples at z>4 to date, confirm less than (HUDF12; Ellis et al. 2013). half of their targeted objects. This is particularly problematic The LBG or “dropout” technique relies on absorption by at z>6 where exotic objects can contaminate the dropout intervening neutral hydrogen below the Lyman limit at 912 Å selection (Capak et al. 2011). and Lyα at 1216 Å to create a strong spectral break that differ- A growing body of evidence suggests extreme line emitters entiates high- and low-redshift galaxies using only broadband and unusual evolved galaxies at z ∼ 2 are a contaminant in photometry. This technique was first introduced in the 1980s z>7 LBG selections (Atek et al. 2011; Capak et al. 2011; and 1990s (Cowie 1988; Steidel et al. 1996) and was broadly Hayes et al. 2012). In published spectroscopic studies of z>7 adopted as the main technique for finding candidate distant galaxies, the vast majority of results are null, with a large fraction sources once it was shown to be effective spectroscopically of detected objects placed at z<3 (Vanzella et al. 2011; Capak on large samples at z ∼ 3–4 (Steidel et al. 1999, 2002) and et al. 2011; Caruana et al. 2012; Hayes et al. 2012; Ono et al. then deployed at ever higher redshifts (Iwata et al. 2003; Ouchi 2012; Bunker et al. 2013). The null results combined with the et al. 2004; Bouwens et al. 2007). The current frontier is to use type of contamination is worrying because it is from a poorly the near-infrared WFC3/IR Hubble Ultra-Deep Field data for understood population of objects and so is difficult to include “J- and H-dropout” galaxies at redshifts z ∼ 8–12 and has led in the simulations required to quantify LBG selection criteria to a handful of tentative detections, most notably the source (Capak et al. 2011). In the near term, this highlights the need UDFj-39546284, a candidate z ∼ 11.9 galaxy (Ellis et al. 2013) for deep spectroscopic studies at 6 8 and faint fluxes needed to understand these ∗ The data presented herein were obtained at the W. M. Keck Observatory, populations. which is operated as a scientific partnership among the California Institute of UDFj-39546284 has been well studied by many authors and Technology, the University of California, and the National Aeronautics and is only detected in the F160W H-band filter with HST WFC3-IR, Space Administration. The Observatory was made possible by the generous financial support of the W. M. Keck Foundation. even though the deepest data currently available exists at both

1 The Astrophysical Journal Letters, 773:L14 (5pp), 2013 August 10 Capak et al.

UDFj-39546284 CDFS21724 CDFS21970 CDFS23549

A WFC3/IR - F160W B ACS - z850 C ACS - z850 D ACS - z850

A z = 11.9? B z = 1.2 C z = 1.3 D z = 1.2

Figure 1. Image cutouts and the two-dimensional MOSFIRE spectra around UDFj-39546284 (A) and three other bright objects are shown. The MOSFIRE slit positions are marked in yellow and the objects highlighted with a cyan circle. Analysis of the bright compact-object spectra CDFS21724 (B) indicates we are losing no more than 10% of the flux due to mask misalignment. The strong line visible in CDFS21970 (C) is Hα at z = 1.3089. A summary of these and other objects in the mask is giveninTable1. (A color version of this figure is available in the online journal.) bluer and redder wavelengths. It was first reported by Bouwens UDFj-39546284, we give a summary of the results from other et al. (2011a), who originally claimed it was at z ∼ 10. Ellis objects on our slit mask in Table 1. et al. (2013) recently improved the depth of the HUDF F160W We adopt a cosmological model with ΩΛ = 0.7, ΩM = 0.3, and the F105W images by 0.2 and 0.5 mag, respectively, and and h = 0.7 and magnitudes in the AB system. added deep F140W imaging which overlaps the blue half of F160W. This new data (HUDF12) favors UDFj-39546284 being 2. DATA at z = 11.9, but still allows the possibility that it could be a strong line emitter at 2