The Suzaku PV observation of the Planetary Nebula BD+30°3639
description
Transcript of The Suzaku PV observation of the Planetary Nebula BD+30°3639
The Suzaku PV observation of the Planetary Nebula BD+30°363
9Mio MURASHIMA, Motohide KOKUBUN (U-Tokyo),
Hiroshi MURAKAMI (ISAS), Kiyoshi HAYASHIDA (Osaka-U), Kyoko MATSUSHITA (Tokyo-U.Sci), Jun'ichi KOTOKU (Titech),
Keith ARNAUD, Kenji HAMAGUCHI (GSFC),Kazuo MAKISHIMMA (U-Tokyo/RIKEN)
Mostly based on the PhD Thesis by Mio Murashima“X-ray Study of Planetary Nebulae”
Nov. 2005: Approved by SWG to use the Suzaku data, on condition that the results be published by the end of March.
Dec. 20, 2005: Thesis submitted to Dept. Astronomy, U.Tokyo.
Jan. 26, 2006: Thesis defense completed successfully. Feb. 14, 2006: Thesis accepted. March, 2006: Degree to be awarded.
To be submitted to ApJL as Murashima et al.
1. Contributed to R&D and pre-launch tests of the HXD.2. Explored X-ray emission from pl. nebulae (with Dr. Kokubun).3. Proposed and planned the PV observation of BD+30 3639 (acce
pted on Sept. 10, 2005).4. Contributed to quantification of the XIS low-E QE change (on xoo
ps).5. Lead data analysis, and discovered a very high C/O ratio.6. Accomplished a fine PhD thesis in only half a year after the laun
ch (perhaps the quickest record among Japanese astro-satellites).
7. Her results selected as one of three press-release topics at the JAS meeting in March 2006.
8. Presently, serving as a duty scientist at USC.
1. Qualification of Mio Murashima(Guideline set by the Steering Comm. on Jan. 13, 2006)
Feb.20, 2006 Suzaku SWG 3/25
BD+30°3639 / HD184738 / V1966 Cyg One of the most well studied planetary nebulae (PNe) (α,δ)=(19 34 45.23, +30 30 58.9) ; (l,b)= (64.79, +5.
02) Distance 1.3±0.2 kpc NH ~ 1e21 cm-2
2. The target
color: H-alphacontours: X-ray
2.5 arcsec The X-ray brightest PN. X-rays are emitted from inside
the optical shell. Similar shapes in C- and O-bands.
(2-1) Basic information
0.3–0.5 keV 0.5–0.7 keV
Feb.20, 2006 Suzaku SWG 4/25
(2-2) Previous X-ray spectroscopy
Newton
Chandra
XIS-BI XIS-FI
Responses to 0.37 keV monoenergetic X-rays
ASCAArnaud et al. (1996)
ChandraKastner et al. (2000)Mannes et al. (2003)
0.5 1.0 1.5 keV
Ne (H)
C N O
C N O Ne
ASCA354 <35 <3
10.5 +7.7-4.5
Chan-dra 354 9 4.2
19.3 ±1.4
Extreme abundances suggested
assumed
Feb.20, 2006 Suzaku SWG 5/25
3. The Suzaku Observation PV-phase observation on 2005 September 21-22 Exposure 34.3 ks 0.033 c/s/XIS-FI, 0.089 c/s/XIS1
SDSS optical imageGray scale: [OIII]Contours: X-rays(ROSAT)Mavromatakis et al. (2002)
SNR G65.2+5.7
18 arcmin 1 deg.
XIS field of view
Feb.20, 2006 Suzaku SWG 6/25
source region bkgd region
5. The XIS Images
XIS-0 (0.3-0.7 keV) XIS-1 (0.3–0.7 keV)
XIS-1: 0.3–0.5 keV
XIS-1: 0.5–0.7 keV
Feb.20, 2006 Suzaku SWG 7/25
6. Analysis of the XIS Spectra(6-1) Raw XIS-1 spectra
kT=0.2 keV 1 solar
[H-like C] / [He-like O] line ratio
Decreasing relative QE
× 0.75
XIS-BI response (initial) × 0.5
Inter-Stellar NH (1e21)
× 0.4
~0.3
Expected count ratio : ~0.02
Excess NH (1e21) × 0.4
The observed C lines are much stronger!
source region,background region
H-like O-Kα0.65 keV
He-like O-Kα0.56 keV
C-Kβ (0.44 keV)
H-like C-Kα (0.37 keV)
He-like Ne-Kα0.91 keV
Feb.20, 2006 Suzaku SWG 8/25
XIS-1XIS-023
Before correction
XIS gain: self-calibrated (2-3 %) using lines in the source spectraLow-E QE: calibrated using RXJ 1856.5-3754 (2005 Oct. 24).
XIS-1XIS-023
After correction
Energy dependence of excess absorption ⇒ modeled so that the XIS spectra of RXJ 1856 can be reproduced by the XMM-determined blackbody model.
Time evolution of the excess absorption ⇒ modeled so that the Chandra flux of BD+30 3639 agrees with that with Suzaku.
QE decrease from the launch at 0.37 keV is ~54% on Sept. 21, and ~30% on Oct. 24.
(6-2) The XIS low-energy CAL (on xoops)
XMM model
Feb.20, 2006 Suzaku SWG 9/25
APEC, NH=1x1021 cm-2
1.5
CVI OVIIOVIII NeIX
0.3 0.5 1.0 2.0
Energy (keV)
(6-3) Simple model fits to the XIS-1 spectrum
Fixed at 1 solar abundance⇒ kT ~ 0.24 keV (χ2/dof = 648/79)
0.3 0.5 1.0 2.0
Energy (keV) Fiexed at solar ratios⇒ kT ~ 0.23 keV, ab.~ 0.01 solar (χ2/dof = 296/78)The spectrum cannot be reproduced by isother
mal solar-ratio IE models.
Feb.20, 2006 Suzaku SWG 10/25
(6-4) 1T analysis vAPEC, ACIS/XIS-1/XIS-023 joint
0.3 0.5 1.0 1.5
Energy (keV)
ACISXIS-1
XIS-023
Absolute abundances are highly uncertain because;- He/H ratio can be non-unity- Metals themselves emit continuum
NH = (2.1+0.2–0.4) e21
kT = 0.19 ± 0.01 keVC = 19+43 -11
N = 0.67+0.25 -0.20 O = 0.20+0.03 -0.02
Ne = 1.1+0.2 -0.1
Fe < 0.07Others = 0 fixedχ2/dof = 312/229
Feb.20, 2006 Suzaku SWG 11/25
Confidence contours68%, 90%, 99%
Abund. (solar)
90% range
C/O 95 75-110
N/O 3.3 1.0-5.0
Ne/O 5.5 4.8-7.3
Abundance ratios are well dtermined.
C/O N/O
Ne/O
(6-5) 1T abundance ratios
Lower-kT plasmas would emit stronger C-lines. What if there are mutiple kT’s (but with the same abundances)? Test 2T-model fits (free kT’s and free norms), but asuuming the C/O ratio.
0
10
20
30
40
50
60
10 100
del
ta-c
his
qu
are
assumed C/O ratio10 20 30 50 100 200
Assumed C/O ratio
(6-6) 2T analysis
010203040506010100delta-chisquareassumed C/O ratio
1T: C/O = 75~110
0.3 0.5 1.0 2.0
Energy (keV)
2T: C/O = 35~120
ACIS
XIS-1
XIS-023
ACIS
XIS-1
XIS-023
The high C/O ratio remains unchanged even considering muti-kT condition.
Feb.20, 2006 Suzaku SWG 13/25
7. Interpretation
Fast wid( 700 km/s )
Central star
Hot gas
Shell (nebula)
Contact discontinuity
Inner shock
Outer shock
Interpretable in temrs of Interacting Stellar Wind model (Kwok et al. 1978; Volk & Kwok 1985)
C/O ~ 60 (solar units)
14N(α,γ)18F(β+ν)18O(α,γ)22Ne
Typical in the He-burning layer; competition between 3α reaction and 12C(α,γ)16O (Suda et al. 2004)
H-rich env.
Masss loss
conduciton
convection
He-burning shell
X-ray emission (size, kT,,..)
Lx~ ~ 1.4×1033 erg s-1
only ~0.1% of kinetic luminosity supplied by the winds
Mass, density, and ionization equilibrium of the plasmaMx~4×10-4 M◎ ; ne~100 cm-3 ; t~103 yr ; ne t~ 3×1012
N/O ~ 3 (solar units)
Ne/O ~ 7 (solar units)
Remainder of CNO cycle
We have successfully detected the He-burning products at the final evolutionary stage of a low-mass star!