The Astrophysical Journal, 780:184 (6pp), 2014 January 10

doi:10.1088/0004-637X/780/2/184

C 2014. The American Astronomical Society. All rights reserved. Printed in the U.S.A.

SN 2005ip: A LUMINOUS TYPE IIn SUPERNOVA EMERGING FROM A DENSE CIRCUMSTELLARMEDIUM AS REVEALED BY X-RAY OBSERVATIONS
Satoru Katsuda1 , Keiichi Maeda2,3 , Takaya Nozawa3 , David Pooley4,5 , and Stefan Immler6,7,8
1 RIKEN (The Institute of Physical and Chemical Research) Nishina Center, 2-1 Hirosawa, Wako, Saitama 351-0198, Japan
2 Department of Astronomy, Kyoto University, Kitashirakawa-Oiwake-cho, Sakyo-ku, Kyoto 606-8502, Japan

3 Kavli Institute for the Physics and Mathematics of the Universe (WPI), University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba 277-8583, Japan
4 Department of Physics, Sam Houston State University, Huntsville, TX 77341-2267, USA
5 Eureka Scientific, Inc., 2452 Delmer Street, Suite 100, Oakland, CA 94602, USA
6 Astrophysics Science Division, NASA Goddard Space Flight Center, Greenbelt, MD 2077, USA
7 Center for Research and Exploration in Space Science and Technology, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA
8 Department of Astronomy, University of Maryland, College Park, MD 20742, USA

Received 2013 September 15; accepted 2013 November 27; published 2013 December 23

ABSTRACTWe report on the X-ray spectral evolution of the nearby Type IIn supernova (SN) 2005ip based on Chandra and Swiftobservations covering ∼1–6 yr after explosion. X-ray spectra in all epochs are well fitted by a thermal emission
model with kT  7 keV. The somewhat high temperature suggests that the X-ray emission mainly arises from thecircumstellar medium (CSM) heated by the forward shock. We find that the spectra taken two to three years after
the explosion are heavily absorbed (NH ∼ 5 × 1022 cm−2 ), but the absorption gradually decreases to the level of the
Galactic absorption (NH ∼ 4 × 1020 cm−2 ) at the final epoch. This indicates that the SN went off in a dense CSM
and that the forward shock has overtaken it. The intrinsic X-ray luminosity stays constant until the final epoch,
when it drops by a factor of ∼2. The intrinsic 0.2–10 keV luminosity during the plateau phase is measured to be∼1.5 × 1041 erg s−1 , ranking SN 2005ip as one of the brightest X-ray SNe. Based on the column density, we derivea lower limit of a mass-loss rate to be Ṁ ∼ 1.5 × 10−2 (Vw /100 km s−1 ) M yr−1 , which roughly agrees withthat inferred from the X-ray luminosity, Ṁ ∼ 2 × 10−2 (Vw /100 km s−1 ) M yr−1 , where Vw is the circumstellar
wind speed. Such a high mass-loss rate suggests that the progenitor star had eruptive mass ejections similar to aluminous blue variable star. The total mass ejected in the eruptive period is estimated to be ∼15 M , indicatingthat the progenitor mass is 25 M .
Key words: circumstellar matter – supernovae: general – supernovae: individual (SN 2005ip) – X-rays: general

Smith et al. (2009b) found clear narrow Hα emission even on
day 1 as well as an unusually rich forest of narrow coronallines at later times and categorized SN 2005ip as SNe IIn.
Fox et al. (2009) reported strong IR emission over more than
900 days post-discovery. Subsequent studies on IR photometricand spectroscopic observations showed two dust components:
one presumably formed in the ejecta or cool dense shell, andthe other, preexisting dust in the surrounding CSM (Fox et al.
2010). X-ray emission from SN 2005ip was detected on 2007
November with Swift (Immler & Pooley 2007). For an assumed
foreground column density of NH = 3.7×1020 cm−2 and athermal spectrum with a plasma temperature of kT = 1 keV,an absorption-corrected (or intrinsic) X-ray luminosity wascalculated to be 1.6±0.3×1040 erg s−1 (Immler & Pooley 2007).
Previous optical and IR observations have led to different
estimates of mass-loss rates for the progenitor of SN 2005ip.
From the Hα intensity together with the CS wind speed(Vw = 100 km s−1 ), Smith et al. (2009b) inferred that Ṁ =(2–4)×10−4 M yr−1 and argued that the progenitor is a RSG
star having eruptive mass ejections like VY Canis Majoris(Smith et al. 2009a). By contrast, a large amount of warm dust
contained in the CS shell led Fox et al. (2010) to suggest avery high mass-loss rate of 0.075–0.38 M yr−1 . Such a massloss rate exclusively favors a LBV-like progenitor rather than a
RSG. Most recently, Moriya et al. (2013) estimated that Ṁ =(1.2–1.4)×10−3 M yr−1 , based on analytical modeling of thebolometric light curve. This result further supports a LBV-likeprogenitor.
In this article, we report on archival X-ray monitoring observations of SN 2005ip from ∼1 yr to ∼6 yr post-discovery,

1. INTRODUCTION
Type IIn supernovae (SNe IIn) are a class of supernovae (SNe)that show prominent narrow H emission lines in their optical
spectra (Schlegel 1990). There is a general consensus that theseSNe occur in dense circumstellar media (CSMs) created by preSN mass losses, and that the narrow H lines are produced byphotoionization of the dense winds irradiated by X-rays from theregion behind the forward shock. However, the detailed natureof SNe IIn is still unclear because well-studied objects are sparse(See Taddia et al. 2013; Kiewe et al. 2012 for recent reviews).
In particular, their progenitors have not yet been established.
One of the most important quantities to constrain the progenitor
is a mass-loss rate, which varies from ∼10−6 –10−5 M yr−1
for red supergiants (RSG), ∼10−5 –10−4 M yr−1 for stripped
Wolf–Rayet stars, and to ∼10−4 –10 M yr−1 for luminous blue
variable (LBV) stars (Kiewe et al. 2012). We can infer massloss rates of progenitor stars by measuring Hα and X-rays from
SNe, as they arise due to interactions between the CSM andejecta. So far, mass-loss rates have been inferred for several
SNe IIn. The values scatter within 10−4 –10 M yr−1 , which isconsistent with massive eruptions seen for LBV stars (Kiewe
et al. 2012; Taddia et al. 2013). Indeed, possible detections of
LBV-like progenitors in pre-SN images were reported for SNe
2005gl and 2009ip (Gal-Yam & Leonard 2009; Mauerhan et al.
2013; Pastorello et al. 2013).
SN 2005ip, discovered on 2005 November 5.163 in the nearby
galaxy NGC 2906 (Boles et al. 2005), was initially designateda normal Type II SN, based on an optical spectrum taken
one day after discovery (Modjaz et al. 2005). Later, however,
1The Astrophysical Journal, 780:184 (6pp), 2014 January 10

Katsuda et al.
Table 1
X-Ray Observations of SN 2005ip

Date
(UT)

Daya

Instrument

Exposure(ks)

Count Rate
(0.5–10 keV)

Hardness Ratio (2.5–10 keV)/(0.5–2.5 keV)

2007 Feb 14.08
2008 Jun 6.46
2008 Nov 16.03
2009 Nov 18.69
2012 Mar 20.02
2012 Mar 22.75
2012 Mar 23.28
2012 Mar 25.62
2012 Mar 31.57

474.92798.091115.871483.53
2335.86
2338.59
2339.122341.462347.41

Swift XRT
Chandra ACIS-S
Swift XRTSwift XRTSwift XRTSwift XRTSwift XRTSwift XRTSwift XRT

8.7
4.7
10.0
6.0
2.4
0.7
0.7
3.3
1.4

4.4 ± 0.7×10−3
22.4 ± 0.2×10−3
10.5 ± 1.0×10−3
10.8 ± 1.3×10−3

7.3 ± 4.3
2.9 ± 0.7
1.9 ± 0.4
1.1 ± 0.3

7.5 ± 0.9×10−3

0.5 ± 0.1

Note. a Days after explosion. For Swift and Chandra data, the PIs are S. Immler and D. Pooley, respectively.

acquired by Swift and Chandra. X-ray emission has been detected in all the observations, exhibiting clear spectral softening
with time. The spectral evolution can be best interpreted by asignificant decline of absorption due to the external CSM. The
X-ray spectra enable us to measure a correct X-ray luminosity
as well as an absorbing column density, allowing for a new estimate of the mass-loss rate, independent of the other previous
ones. Throughout this article, we use a distance to SN 2005ip of35 Mpc (Stritzinger et al. 2012) and adopt 2005 October 27.2
(nine days prior to the discovery) as the time of explosion (Smithet al. 2009b).

Ratio

Normalized counts s−1 keV−1

100

2. OBSERVATIONS AND ANALYSIS

2007−02(×50000)
2008−06(×500)

1
0.1
0.01

2008−11(×100)
2009−11(×10)
2012−03

10−3
10−43
2
1
0
0.5

We analyze several archival Swift and Chandra data for SN
2005ip using the HEAsoft9 and CIAO10 packages. Detailed
information of these observations is summarized in Table 1.
We combine the latest five Swift observations taken in 2012
March into one data file owing to the poor photon statistics
for individual data sets. As a result, we have five observation
epochs, i.e., 2007 February (Swift), 2008 June (Chandra), 2008
November (Swift), 2009 November (Swift), and 2012 March(Swift).
We extract X-ray spectra from a circular region with a radiusof 30 (3 ) for Swift (Chandra) observations and subtract background from the surrounding annular region. The backgroundlevels to the source are quite small: ∼0.3% (Chandra) and ∼3%(Swift). Each spectrum is grouped into bins, so that each bin
contains at least five counts. We have checked that even if we
set the number of the minimum grouping count to be unity,
we obtain the same fit results shown below. The backgroundsubtracted spectra are shown in Figure 1. Clearly, the spectra aresoftening with time, as is evident in hardness ratios summarizedin Table 1.
For spectral modeling, we apply an absorbed (TBabs; Wilms
et al. 2000) thermal emission model in ionization equilibrium,i.e., the apec model (Smith et al. 2001) in XSPEC (Arnaud
1996). We allow the absorbing H column density (NH ), thetemperature (kT ), and normalizations to be free, while we treat
the temperature to be common among all the five spectra inorder to better constrain the parameters. Elemental abundancesin both TBabs and apec models are set to the solar values(Lodders 2003), which are consistent with the metallicity of
the host galaxy NGC 2906 (Stritzinger et al. 2012). To findthe best-fit model, we use maximum likelihood statistics for
9

10

1

2
Energy (keV)

5

Figure 1. X-ray spectra along with the best-fit models (TBabs×apec). All thedata were obtained with the Swift XRT except for the one in 2008 June, which
was taken with the Chandra ACIS-S. The lower panel shows the ratios of thedata to the best-fit model.

a Poisson distribution, the so-called c-statistics
(Cash 1979),
which minimizes the C value defined as C = −2 i (Mi − Di +Di (lnDi − lnMi )) with Mi and Di being the model-predicted
and observed counts in each spectral bin i, respectively. Thismethod is suitable for low-counts spectra that do not allow forthe χ 2 test. Errors on the parameters can be estimated similarly
to the χ 2 method, based on the fact that ΔC = C − Cmin can beused similarly to Δχ 2 (Cash 1979), where Cmin is the minimum
C value obtained at the best-fit. As can be seen in Figure 1, thebest-fit models represent the data well. The best-fit parameters
are summarized in Table 2 (No. 1). We find a clear temporaldecrease in NH . This result is fairly robust from a statisticalpoint of view, as is evident from confidence contours of kT
versus NH in Figure 2, where thick and thin lines correspondto the first and the final epochs, respectively. Judging from thehigh temperature (6.7 keV), most of the X-ray emission likely
arises from a hot plasma in the CSM heated by the forward
shock.
We also fit the data with the same model in different waysthat take account of temperature evolution. First, we assume
that the temperature evolves as T ∝ t −0.2 , as the temperature is
proportional to the square of the shock speed, which decreases
as t −0.1 in a self-similar solution (e.g., Chevalier 1982); the SN
radius, R, is proportional to t (n−3)/(n−s) , where n and s are density
profiles in the ejecta and the CSM, respectively (ρej ∝ r −n andρcsm ∝ r −s ), so that the shock speed, dR/dt, is proportional to

http://heasarc.gsfc.nasa.gov/docs/software/lheasoft/

10 http://asc.harvard.edu/ciao/

2

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Katsuda et al.

Table 2
Spectral Properties from the One-component Model
Parameter
NH(1022 cm−2 )

1
2
3
4

kT (keV)

Observed flux (10−13 erg s−1 cm−2 )Absorption-corrected flux (10−13 erg s−1 cm−2 )2008–11

2009–11

2012–03

3.12+1.44−0.87
3.11+1.36−0.87
3.08+0.69−0.61

1.35+0.89−0.51
1.33+0.88−0.49
1.33+0.45−0.38

0.86+1.04−0.68
0.85+1.05−0.66
0.86+0.55−0.47

0.21(<0.41)

0.21(<0.5)
0.22(<0.46)

Linked to 2007–02

1

68.6(> 6.7)

Linked to 2007–02

81.5(> 8.5)

Linked to 2007–02 with assumed evolution of t −0.2

3172.6(> 109.6)

Linked to 2007–02 with assumed evolution of t −0.53

4135.0(>52.6)

100.0(>26.3)

87.8(>26.3)

7.1+39.7−3.2

3.0+1.5−0.9

1

5.69+1.67−1.40

6.00+1.03−0.93

8.11+1.42−1.27

7.80+1.79−1.56

4.46+1.00−0.87

25.73+1.68−1.41

6.0+1.0−0.9

8.17+1.43−1.28

7.84+1.80−1.56

4.48+1.0−0.87

35.73+1.68−1.41

6.04+1.04−0.93

8.17+1.43−1.28

7.80+1.79−1.56

4.40+0.98−0.86

4

3.91+1.14−0.96

4.95+0.85−0.77

8.55+1.50−1.34

6.91+1.59−1.38

3.46+0.79−0.69

19.65+2.83−2.37

8.94+1.54−1.38

10.86+1.90−1.71

9.91+2.28−1.98

5.03+1.12−0.98

28.93+1.54−1.38
8.93+1.54−1.38
6.63+1.14−1.02

10.85+1.90−1.71
10.85+1.90−1.70
11.65+2.04−1.83

9.91+2.28−1.97
9.91+2.28−1.98
11.20+2.57−2.23

5.04+1.12−0.98

4

9.61+2.82−2.36
9.61+2.82−2.36
5.34+1.55−1.31

11.43+0.42−0.35

1.32+0.23−0.20

1.61+0.28−0.25

1.47+0.34−0.29

0.74+0.17−0.15

2

1.42+0.42−0.35
1.42+0.42−0.35
0.79+0.23−0.19

1.32+0.23−0.20
1.32+0.23−0.20
0.98+0.17−0.15

1.60+0.28−0.25
1.60+0.28−0.25
1.72+0.30−0.27

1.47+0.34−0.29
1.47+0.34−0.29
1.66+0.38−0.33

0.75+0.17
−0.15

3
4
C/dof

2008–06

5.38+3.02−2.18
5.35+3.09−2.16
5.33+2.10−1.55
1.62+0.25
−0.22

2

3
La (1041 erg s−1 )2007–02

12
3
4

5.01+1.12−0.98
7.47+1.71−1.49

0.74+0.17−0.14
1.10+0.25−0.22

56.1/6256.1/6256.0/6293.1/62

Note. a Calculated in 0.2–10 keV at a distance of 35 Mpc.

nonequilibration may be applicable for SN 2005ip, in which
X-rays are likely due mainly to an adiabatic forward shock(e.g., Fransson et al. 1996). To this end, we follow Nymark
et al. (2009) who showed that the electron temperature is−4/3
proportional to Vsh t −2/3 for Coulomb heating (Equation (10)in Nymark et al. 2009). Using the relation Vsh ∝ t −0.1 , we obtainTe ∝ t −0.53 . By fitting the five spectra with this temperature
evolution taken account, we obtain the results shown in Table 2(No. 3), which are consistent (except for the temperature) withthe other cases. Finally, we allow kT to vary freely among thefive spectra, while we let NH be a common free parameter. This
strategy fails to fit the data (C value of 93.1; the fit results arepresented in Table 2 No. 4), indicating that the spectral softening
is not simply due to cooling. Therefore, we are confident that the
spectral softening is mainly due to evolution in the absorption.
While non-detection of X-ray emission from reverse-shocked
ejecta is not so peculiar for SNe IIn (cf. Chandra et al.
2012a, 2012b), for completeness we also apply an absorbed
two-component apec model, in which the two components
are supposed to represent the reverse-shocked ejecta and theforward-shocked CSM (e.g., Fransson et al. 1996; but see also
Nymark et al. (2006, 2009) who presented a more detailed
modeling especially focusing on X-ray emission from the ejecta
by coupling a spectral code to a hydrodynamical code). Since
it is difficult to determine a number of spectral-fit parameterswith the relatively poor photon statistics, we initially linktemperatures and normalizations of the low-temperature (ejecta)
component among the five spectra, according to a theoretical

50

kT (keV)20
1052012 March

2007 February

2
1

0.1

1
NH (1022 cm−2)

10

Figure 2. Confidence contours of kT vs. NH for which black, red, and blue
represent confidence levels of ΔC = 2.3, 4.6, and 9.2 (corresponding to 68%,
90%, and 99% for two interesting parameters). Thin and thick lines correspondto the results taken in 2007 February and 2012 March, respectively.

t (n−3)/(n−s)−1 , resulting in dR/dt ∼ t −0.1 for typical density
profiles of core-collapse SNe, n ∼ 10 and s ∼ 2 (e.g.,
Chevalier & Fransson 2003). The fit results listed in Table 2 (No.
2) are found to be consistent (except for the temperature) withthose of the No. 1 case. Second, we take into account effects oftemperature non-equilibration between electrons and ions due toslow equilibration through Coulomb heating. The temperature
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Katsuda et al.

Table 3Spectral Properties from the Two-component Model
Parameter
NH(1022 cm−2 )12
3

kTl (keV)

Lal (1041 erg s−1 )Lah (1041 erg s−1 )C/dof

2008–06

2008–11

2009–11

2012–03

7.24+3.27−2.19
2.94+0.79−0.52
3.26+0.98−0.90

5.06+1.29−1.64

3.14+1.04−1.34

2.78+1.15−1.05

2.32+0.91−1.29

3

0.22+0.70−0.07
0.22+0.07−0.07
0.30+0.15
−0.12

1
2

kTh (keV)

2007–02

Linked to 2007–02
Linked to 2007–02
Linked to 2007–02
Linked to 2007–02
Linked to 2007–02 with assumed evolution of t −0.2

1

9.19(> 4.63)

Linked to 2007–02

2
3

88(> 36)102.3(> 9.0)

Linked to 2007–02
Linked to 2007–02 with assumed evolution of t −0.2

135.0+10.4−9.3

20(<5.2)0(<5.8)35.5+20.7−16.6

49.3+34.7−27.0

74.9+26.9−22.6

30(<1.46)0(<2.63)

19.6+11.3−9.2

35.5+24.1−19.0

82.3+29.7−24.9

1

1.73+0.52−0.43

1.58+0.29−0.26

1.90+0.38−0.33

1.81+0.48−0.41

0.91+0.26−0.22

2

1.04+0.30−0.25

1.29+0.22−0.20

1.84+0.38−0.34

1.80+0.52−0.45

0.80+0.29−0.23

3

1.04+0.30−0.25

1.29+0.22−0.20

1.82+0.38−0.34

1.76+0.53−0.45

0.80+0.29−0.24

Linked to 2007–02

12
3

51.6/60
58.4/60
58.0/60

Note. a Calculated in 0.2–10 keV at a distance of 35 Mpc.

expectation by Fransson et al. (1996). The best-fit parameters are
summarized in Table 3 (No. 1). The two-component model givesa slightly better fit than the one-component model: ΔC ∼ −4.5
for two extra free parameters (i.e., kT and normalization of
the ejecta component). The slight improvement in the fit is notstrongly significant, yielding a confidence level of 92% basedon the F-test; we assume that the C values can be treated as
χ 2 values and performed F-test using the fit-statistics listed
in Table 3 (No. 1). It should be noted that, as can be seenfrom the best-fit parameters listed in Tables 2 and 3, the basic
best-fit parameters, i.e., NH , kT , and normalization of the hightemperature component, derived by the one- and two-component
models, roughly agree with each other.
Next, to see if the spectral softening can be explained by
the increasing low-temperature component rather than the NH
evolution, we fit the data with the two-component model in adifferent fitting strategy, in which we link the NH parameter
among the five epochs while we let the normalization of the
low-temperature component vary freely for individual epochs.
The best-fit parameters are listed in Table 3 (No. 2), fromwhich we find that the model gives a reasonable but slightly
worse fit than previous fittings. Moreover, the intrinsic X-ray
luminosity for the low-temperature component is derived tobe ∼1043 erg s−1 at the final epoch (∼6 yr after explosion).
This may be unrealistically bright especially for reverse-shock
emission; it is two orders of magnitude brighter than everrecorded at this late phase (Dwarkadas 2012). The intrinsic
X-ray luminosity for the cooler component would be evengreater than the value we obtained. This is because there isalmost always a radiative cooling shell between a reverse shock
and a forward shock, and the cooling shell will provide an extra
absorption of the soft X-rays coming from the reverse shock.
We have also checked that temperature evolution of T ∝ t −0.2
for both of the two components (as is expected in a reasonableself-similar solution described above) does not affect the fit
result (see the results No. 3 in Table 3). Accordingly, in thefollowing discussion, we basically adopt the results from theone-component model in Table 2.
3. DISCUSSION
The archival X-ray data of SN 2005ip revealed clear spectral
softening with time. Our spectral analysis showed that the
spectral softening is best explained by declining evolution of
the absorbing column density; NH ∼ 5 × 1022 cm−2 at thefirst epoch (∼1 yr after explosion) gradually goes down to beconsistent with the Galactic absorption, NH ∼ 4 × 1020 cm−2 ,at the final epoch (∼6 yr after explosion). On the other hand,
we find that the intrinsic X-ray luminosity stays constant until
the final epoch when it drops by a factor of ∼2 (see Table 2).
These results suggest that the SN was initially embedded in adense CSM, but the forward shock has penetrated the dense
CSM region at a certain time between 2009 November (fourth
epoch) and 2012 March (final epoch).
SN 2005ip is the second example in which we see possible
evolution in absorption by the external CSM after SN 2010jl(Chandra et al. 2012a). By contrast, there are several SNethat do not show significant NH evolution. These include SNe
2006jd (Chandra et al. 2012b), 1987A (Dewey et al. 2012),
1998S (Pooley et al. 2002), and 1999em (Pooley et al. 2002).
SN 1993J would be intermediate between the two types, i.e.,evolving absorption and constant absorption; Chandra et al.
(2009) noted possible NH evolution that may be caused by a coolshell between the forward shock and the reverse shock. A simple
speculation to explain the difference between the two types isthat the CSM is formed as a disk-like shell due to equatorial
mass ejections. For this geometry, CSM absorption strongly
depends on the inclination angle of the disk in the sense that
4

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Katsuda et al.

than the original value of 1.6 ± 0.3 × 1040 erg s−1 reported by
Immler & Pooley (2007). The revised luminosity places SN
2005ip as one of the brightest X-ray SNe among others being
SNe 2001em (Pooley & Lewin 2004), 2005kd (Immler et al.
2007; Pooley et al. 2007; Dwarkadas 2012), 2006jd (Chandra
et al. 2012b) and 2010jl (Chandra et al. 2012a). By integrating the equation from 0.2 keV to 10 keV, it can be rewritten as 2.77×1036 (1−s)−1 (Ṁ−4 /Vw2 )2 (Vsh /104 km s−1 )3−2s(t/11.57 days)3−2s erg s−1 for T = 6 × 108 K, where Ṁ−4 =(Ṁ/10−4 M yr−1 ) and Vw2 = (Vw /100 km s−1 ). The intrinsic
luminosity during the plateau phase, first–fourth epochs when
we can reasonably assume that s = 1.5 from 3 − 2s = 0,
gives a mass-loss rate of Ṁ = (2–2.5) × 10−2 Vw2 M yr−1 ,which is in agreement with the lower limit of the NH -based Ṁ
estimated above. Therefore, our X-ray-based mass–loss rates
are significantly higher than those inferred from the Hα luminosity ((2–4) × 10−4 M yr−1 ; Smith et al. 2009b) andthe bolometric light curve ((1.2–1.4)×10−3 M yr−1 ; Moriya
et al. 2013), but are close to the estimate from IR observations((7.5–38)×10−2 M yr−1 ; Fox et al. 2010). Our result supports
eruptive mass ejections as seen in LBV stars.
Interestingly, the evolution in the X-ray luminosity, which isconstant for the initial ∼4 yr and drops between the fourth epoch
and the final epoch, is similar to that of Hα emission (Stritzinger
et al. 2012). This similarity confirms a general picture that latetime Hα emission arises by reprocessing of X-ray emission. The
decreasing X-ray and Hα luminosities at the final epoch lead us
to suggest that the forward shock broke out a dense CSM region(i.e., entered into a steeper density-profile region) at some point
between the fourth epoch and the final epoch. The distance thatthe forward shock (Vsh = 14,000 km s−1 ; Stritzinger et al.
2012) traveled for ∼5 yr, a time span for the forward shock
to be within the dense CSM region, is calculated to be ∼2.2 ×
1017 cm. This would be the spatial extent of the erupted material.
Dividing the distance by the wind speed (Vw ∼ 100 km s−1 ;Smith et al. 2009b), we can estimate the duration of the eruption
to be ∼700 yr (t/5 yr) (Vsh /14,000 km s−1 ) (Vw /100 km s−1 )−1 .
Then, the total mass ejected in this eruptive period is calculated
to be ∼14 (t/700 yr) (Ṁ/0.02 M yr−1 ) M , indicating that theprogenitor is quite massive (25 M ).
As mentioned above, the constant evolution in the X-ray
luminosity (the plateau phase) requires that s = 1.5 in theframework of general SN evolutionary models (e.g., Fransson
et al. 1996). At a glance, this appears to conflict with the NH based s-value ranging from ∼2 to ∼3. However, we should notethat such SN evolutionary models assume spherically symmetric
CSM structures, which may not be valid for SN 2005ip. For
example, the CSM geometry may be a disk-like shell, as
discussed earlier in this section. In this case, the constant X-ray
luminosity is expected for a density profile of s = 2 instead ofs = 1.5 (Stritzinger et al. 2012). On the other hand, evolution in
NH would be similar to that expected for a spherically symmetric
CSM, if the disk is edge on. Therefore, a disk-like CSM structurewith s = 2 would explain evolution in both NH and luminosity
simultaneously without invoking exotic conditions.
There still remains the possibility that the CSM structure of
SN 2005ip is spherically symmetric. In this case, evolution in
NH and the luminosity leads to discrepant density profiles, which
seems to be a real problem. By noting that absorbing material
decreases more rapidly than what is expected from the luminosity evolution, one may think of a possibility that the external
cool CSM is fully ionized by strong X-ray emission behindthe forward shock. We thus calculate the ionization parameter,

Forward shock position (cm)

0.1

500

2012 Mar.

2008 Jun.

2007 Feb.

NH (1022 cm−2)10

12x10172009 Nov.

1017

2008 Nov.

5x1016

1000
Days since explosion

2000

Figure 3. NH as a function of time after explosion along with two model curves
for either Ṁ = 1.5×10−2 Vw2 M yr−1 and s = 2 (solid) or Ṁ = 4Vw2 M yr−1
and s = 3 (dotted). The upper label shows an approximate forward shock position
based on the shock speed of Vsh = 14000 (t/1 yr)−0.1 km s−1 (cf. Stritzinger
et al. 2012).

strong absorption and its declining evolution is expected for anedge-on view (which would be responsible for SNe 2005ip and2010jl), while little absorption for a face-on view (for others). An
alternative scenario is that the class of SNe without evolution
in absorption has somewhat peculiar CSM geometries (while
the SNe class with strongly evolving absorption has either arelatively spherically symmetric CSM or an edge-on disk). For
example, Chandra et al. (2012b) proposed for SN 2006jd thata low-column-density region is present in one direction (i.e., acylindrical shape toward the observer) while strong interactions
are taking place over much of the rest of the solid angle. In order
to constrain the CSM structures, it is essential to increase SNesamples that allow for absorption measurements with Chandra,
XMM-Newton, and Swift.
We estimate mass-loss rates of the progenitor star of
SN 2005ip in two different ways. One is based on the column density, which is a robust measure in X-ray spectroscopy.
We compare the measured column density with a theoretical
calculation, i.e., Equation (4.1) in Fransson et al. (1996). As
shown in Figure 3, the data are well represented by eitherṀ = 1.5 × 10−2 M yr−1 and s = 2 (solid) or Ṁ = 4 M yr−1
and s = 3 (dotted) with reasonable wind and shock speeds(Vw =100 km s−1 and Vsh = 14,000 km s−1 ; Smith et al.
2009b). Here, we should not put much weight on the last-epochdata, since the forward shock has likely penetrated the dense
CSM region (resulting in a change in the s value) between the
fourth epoch and the final epoch. From an astrophysical point ofview, the latter case is not realistic; Ṁ = 4 M yr−1 is too highand s = 3 is too steep especially for Type IIn SNe, which areembedded in usually dense environments. Therefore, we favor
the former case with the mass-loss rate of ∼0.015 M yr−1 ands = 2.
The other way is based on the X-ray luminosity, as is often used to estimate wind parameters in conjunction with
Equation (3.11) in Fransson et al. (1996), which describes
X-ray emission from the shocked CSM. The much larger absorption than that (NH = 3.7 × 1020 cm−2 ) assumed by Immler &Pooley (2007) boosts up the intrinsic 0.2–10 keV luminosity tobe ∼1.5×1041 erg s−1 , which is by an order of magnitude higher
5

The Astrophysical Journal, 780:184 (6pp), 2014 January 10

Katsuda et al.

ξ = L/nr 2 . For the X-ray luminosity of L =1.5×1041 erg s−1 ,a lower limit density of n = 3 × 105 cm−3 (Smith et al. 2009b),
and a radius of r = 1017 cm, which is a distance for a shock of
Vsh = 104 km s−1 to propagate during ∼1000 days, we obtainthat ξ  50. This is not sufficient to fully ionize C, N, and O,
which are main absorbers of X-rays (cf. Chandra et al. 2012b).
Therefore, photoionization is unable to explain the discrepancy
in the derived density profile. The other possibility is to introduce clumpiness in the CSM wind, in which shocked clumps
sustain the high X-ray luminosity (Chugai & Danziger 1994),
while a covering fraction of dense clumps is so small that they do
not noticeably affect dynamics of the system and the CSM absorption. In this model, a flatter distribution of clumps than thatof a global CSM density profile can qualitatively explain boththe constant intrinsic luminosity and the rapidly declining NH .
Given that the presence of small clumps is strongly suggestedby the late-time development of intermediate-width componentof Hα (Smith et al. 2009b), this scenario may be at work. The
clumps will produce significant soft X-ray emission, which is
not required in the one-component model but is seen in the twocomponent model (cf. Table 3). Therefore, the clump scenario
cannot be excluded at the expense of a remarkably strong X-ray
emission from the reverse shock.

is supported by World Premier International Research Center
Initiative (WPI Initiative), MEXT, Japan. S.K. and K.M. aresupported by Grant-in-Aid for Scientific Research for Young
Scientists (25800119, 23740141). T.N. has been supported by
the Grant-in-Aid for Scientific Research of the Japan Society
for the Promotion of Science (22684004, 23224004).
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4. CONCLUSION
X-ray spectra from SN 2005ip showed strong spectral softening with time, which we interpret to be due to declining inabsorption by the external CSM. With the absorption taken into
account, SN 2005ip turned out to be one of the brightest SNe IIn
in the X-ray regime. Both the column density and the high luminosity require a mass-loss rate of at least ∼1×10−2 Vw2 M yr−1 .
Such a high mass-loss rate suggests eruptive mass ejections by a
LBV-like progenitor. Based on evolution in the X-ray luminosity and NH , we speculate that the CSM structure of SN 2005ip
is a disk-like equatorial shell. Alternatively, a relatively spherical wind with small clumps could be another possibility at theexpense of a remarkably bright reverse shock X-ray emission.
We are grateful to Drs. Jonathan Mackey and Takashi Moriya
for kindly pointing out a correction to Figure 3 prior to
publication. S.K. is supported by the Special Postdoctoral
Researchers Program in RIKEN. The work by K.M. and T.N.

6