X-ray study of extended emission around M86 observed with Suzaku
Ukyo Hishi, Ryuichi Fujimoto, Misato Kotake, Hiromasa Ito, Keigo Tanaka, Yu Kai, Yuya Kinoshita
aa r X i v : . [ a s t r o - ph . H E ] F e b Publ. Astron. Soc. Japan (2017) 00(0), 1–17doi: 10.1093/pasj/xxx000 X-ray study of extended emission around M86observed with Suzaku
Ukyo H
ISHI , Ryuichi F
UJIMOTO , Misato K
OTAKE , Hiromasa I TO , KeigoT ANAKA , Yu K AI , Yuya K INOSHITA
Faculty of Mathematics and Physics, Kanazawa University, Kakuma-machi, Kanazawa,Ishikawa 920-1192 ∗ E-mail: [email protected], [email protected]
Received 2017 January 10; Accepted 2017 February 15
Abstract
We analyzed the Suzaku data of M86 and its adjacent regions to study the extended emissionaround it. The M86 core, the plume, and the tail extending toward the northwest were clearlydetected, as well as the extended halo around them. From the position angle ∼ ◦ to ∼ ◦ , the surface brightness distribution of the core and the extended halo was representedrelatively well with a single β -model of β ∼ . up to 15 ′ –20 ′ . The X-ray spectra of the core wasrepresented with a two-temperature model of kT ∼ . keV and ∼ . keV. The temperaturesof the core and the halo have a positive gradient in the center, and reach the maximum of kT ∼ . keV at r ∼ ′ , indicating that the halo gas is located in a larger scale potential structurethan that of the galaxy. The temperatures of the plume and the tail were . ± . keV and . ± . keV. We succeeded in determining the abundances of α -element separately, forthe core, the plume, the tail, and the halo for the first time. Abundance ratios with respectto Fe were consistent with the solar ratios everywhere, except for Ne. The abundance of Fewas ∼ . in the core and in the plume, while that in the tail was ∼ . , but the differencewas not significant considering the uncertainties of the ICM. The abundance of the halo wasalmost the same up to r ∼ ′ , and then it becomes significantly smaller (0.2–0.3) at r > ∼ ′ ,indicating the gas with low metal abundance still remains in the outer halo. From the surfacebrightness distribution, we estimated the gas mass ( ∼ × M ⊙ ) and the dynamical mass( ∼ × M ⊙ ) in r < kpc. The gas mass to the dynamical mass ratio was − – − ,suggesting a significant fraction of the halo gas has been stripped. Key words:
Galaxies: individual: M86 — Galaxies: ISM — X-rays: galaxies: clusters
M86 (NGC 4406) is a bright elliptical galaxy in the Virgo clus-ter, located about 1 . ◦
26, or about 350 kpc in projection, fromthe Virgo cluster center M87. Its redshift is z = − . ± . (Cappellari et al. 2011), i.e., it is approaching us withthe line-of-sight velocity of ± km s − . On the other hand,the redshift of M87 is z = 0 . ± . , and it is goingaway from us with the line-of-sight velocity of ± km s − .It is also reported that M86 is only about 1 Mpc more distance than M87 (Mei et al. 2007). Therefore, M86 is likely moving inthe Virgo cluster with a relative line-of-sight velocity of about1500 km s − with respect to the intracluster medium (ICM).This is much larger than the velocity dispersion of galaxies inthe Virgo cluster ( ∼ km s − ) (Binggeli 1999), and hence,the direction of motion is considered close to the line-of-sightdirection. Since the sound speed is 730 km s − for the clus-ter ICM of kT = 2 keV, M86 is moving with a Mach numberof > ∼ . M86 is thought to be the dominant member of one c (cid:13) Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0
Table 1.
Datasets used in the analysis.
No. ObsID Object Obs Date Exposure (ks)1 803043010 NGC 4406 (M86) 2009-06-19 1022 808045010 NGC 4438 Tail 2013-12-10 1033 800017010 NGC 4388 2005-12-24 124 of the subgroups within the Virgo cluster (e.g., B¨ohringer etal. 1994; Schindler et al. 1999). Therefore, M86 provides agood opportunity to study the interaction between the interstel-lar medium (ISM) and the ICM as well as the interaction be-tween the subcluster and the ICM.Characteristic features of M86 were reported by various au-thors, especially in the X-ray band, since it is sensitive to thehot ISM in the elliptical galaxies. Forman et al. (1979) discov-ered a plume of soft X-ray emission, which is thought to bestripped from M86 by ram-pressure with the Virgo ICM (seealso White et al. 1991). Using ROSAT data, Rangarajan etal. (1995) showed that the temperatures of the galaxy and theplume are both ∼ . keV. Using Chandra data, Randall et al.(2008) discovered a very long tail toward northwest, of 150 kpcin projection and the true length of > ∼ kpc. They also de-tected a discontinuity of the X-ray surface brightness, whichwas interpreted as the density jump due to the shock. M86 wasalso observed by XMM-Newton (Finoguenov et al. 2004; Ehlertet al. 2013). Ehlert et al. (2013) examined the temperature andabundance distributions in detail, and reported the existence ofcool ( ∼ . keV) gas trailing to the northwest of M86 and, alsoto the east of M86 in the direction of NGC 4438.In this paper, we report the results of our analysis of theSuzaku archival data of M86 and its adjacent regions, to studythe extended emission around M86. We adopt a distance to theVirgo cluster of 16.5 Mpc, which gives a scale of 4.8 kpc per 1 ′ .All error ranges are 90% confidence intervals, and the F -testsignificance level is 1%, unless otherwise stated. We used three datasets of Suzaku version 2.5 products, archivedin Data ARchives and Transmission System (DARTS) atISAS/JAXA. M86 was observed on 2009 June 19. Adjacentpointings aiming at NGC 4438 and NGC 4388 were also used.They are summarized in table 1.
HEASoft 6.15.1 was used for data processing, extractionand analysis. The data were reprocessed using aepipelinev1.1.0 , with
CALDB version 20150105 for the dataset
CALDB version 20140520 for Version 5.0. There were four inde-pendent XIS units (XIS0–3), but XIS2 was inoperative since http://heasarc.gsfc.nasa.gov/docs/suzaku/analysis/abc/ × and × editingmodes were combined per sensor after the reprocessing. Theregions which were illuminated by the calibration sources werediscarded. The net exposure time of each observation is sum-marized in table 1. The average count rate of XIS1 in the 0.5–5 keV band was 2.0, 1.0, and 0.93 counts s − , respectively. Itwas checked that there was no statistically significant variationin the light curves of the cleaned data.Contribution of the particle background (Non-X-rayBackground; NXB) of each XIS unit was estimated using xisnxbgen tool and data taken when the satellite saw the nightside of the Earth stored in the CALDB , by sorting them by the cut-off rigidity values, and properly weighting them by the exposuretime ratio, based on the results by Tawa et al. (2008). The de-tector redistribution matrix files (RMFs) were generated with xisrmfgen , using the appropriate calibration files at the time ofthe observation. On the other hand, responses of the X-ray tele-scopes were implemented into ancillary response files (ARFs),using ray-tracing based generator xissimarfgen (Ishisaki et al.2007). Time- and position-dependent contamination in the opti-cal path of each sensor was also considered by xissimarfgen .When the ARFs were generated, we assumed a uniform sourcedistribution in a circle of 20 ′ radius. A mosaic of the three pointings of the XIS in the 0.8–1.2 keVenergy band is shown in figure 1. The energy range was se-lected to be sensitive to the hot gas of kT ∼ keV. Images inthis energy band were extracted from the event files of XIS0, 1,2, 3. They were rebinned by a factor of 8 (0 . ′ xisnxbgen , and they were subtracted. Then,flat field images were generated, and the vignetting of the X-raytelescopes was corrected by dividing the XIS images with theflat field images. The mosaic was generated, the correspondingexposure map was generated using xisexpmapgen , and the mo-saic was divided by the exposure map. Finally, it was smoothedwith a Gaussian of σ = 0 . ′ (3 rebinned pixels).As clearly seen in figure 1, an extended emission of a char-acteristic shape with two peaks was detected at the location ofM86, together with two more relatively weak sources at the po-sition of NGC 4388 and NGC 4438. The brightest peak is theM86 center. On the north side of it, a large plume of emissionis seen, and an elongated tail extends toward the northwest. Anextended halo of the X-ray emission is seen around the cen-ter of M86, which extends near NGC 4388 and NGC 4438.All these are consistent with the previously reported structureswith high spatial resolution by ROSAT (Rangarajan et al. 1995),
Chandra (Randall et al. 2008) and
XMM-Newton (Ehlert et al. ublications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 dead column extended halo TailPlume M86 centerNGC4388NGC4438 NE Fig. 1. (Left) Background-subtracted, vignetting-corrected mosaic of XIS images around M86 in the 0.8–1.2 keV energy band. Images of XIS0, 1, 3 werecombined, and it was smoothed with a Gaussian of σ corresponding to 0 . ′
42. (Right) X-ray contour map overlaid with an optical image. The optical image wastaken from the Digitized Sky Survey. . ◦ ′ , and they become flatter outside it.In the the southern regions (sectors 6–8), on the other hand,the slopes change at 10 ′ –12 ′ . Then, there is a contribution ofNGC 4388. Outside ∼ ′ (sectors 7 and 8), the surface bright-ness becomes very low. At r > ∼ ′ in the eastern regions and r > ∼ ′ in the southern regions, it appears that the ICM of theVirgo cluster becomes dominant. Note that a factor of 2–3 dif-ference in the ICM flux was reported around the M86 region(Rangarajan et al. 1995). The surface brightness variation ofthese regions is qualitatively consistent with their results.To understand the surface brightness profiles more quantita-tively, we fitted the profiles of sectors 2, 3, 4, and 6, with a β model (Cavaliere & Fusco-Femiano 1976) S ( θ ) = S (cid:26) (cid:16) θθ (cid:17) (cid:27) − β + + constant , (1)where a constant was introduced to represent the ICM and otherbackground and foreground components. The results are sum-marized in table 2 and the best fit models for sectors 2, 3, 4, and 6 are shown in figure 3. Since it was not possible to con-strain the constant for sectors 3 and 4, it was fixed at the averagevalue of sectors 2 and 6. The surface brightness of the Suzakuimage can be represented relatively well with a single β -modelof β ∼ . , up to 15 ′ –20 ′ in the eastern and southern regions. As described in the previous section, the mosaic of the SuzakuXIS images showed structures of the X-ray emission from M86,i.e., emissions from the M86 core, the plume, the tail, and adiffuse emission around them. We defined regions as shownin figure 4, and performed a model fitting to the spectral dataextracted from these regions. As an emission model froman optically-thin thermal plasma in collisional ionization equi-librium, we used the APEC (Astrophysical Plasma EmissionCode) model (Smith et al. 2001), v2.0.2. It was used for thegalaxy hot gas, the ICM, and the galactic foregrounds, i.e., theLocal Hot Bubble (LHB) and the Milky Way Halo (MWH). Thetemperatures of the LHB and MWH were fixed at 0.11 keV(LHB) and 0.3 keV (MWH), respectively. The solar abundances( Z ) by Lodders (2003) were adopted in the fitting, and thoseof LHB and MWH were assumed the same as the solar val-ues ( Z ⊙ ). The Cosmic X-ray Background (CXB) was mod-eled with a power law function of a photon index 1.4 and anormalization corresponding to 9.7 photons s − cm − at 1 keV(Revnivtsev et al. 2005). The galactic hydrogen column densitywas fixed at . × cm − (Kalberla et al. 2005). Note thata constant ratio between the BI CCD data and the FI CCD data Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0
Sector 10Sector 9Sector 8Sector 5Sector 3 Sector 1 Sector 7Sector 6Sector 4Sector 2 (cid:0) . × (cid:0) × (cid:0) × (cid:0) . S u r f a c e b r i gh t ne ss ( c oun t s s (cid:0) a r c m i n (cid:0) ) Angular distance from M86 center (arcmin)
Fig. 2. (Left) Definition of sectors 1–10, to examine the surface brightness profiles. Dotted circles show the distance from the M86 center, 5, 10, 15, 20, 25arcmin, respectively. (Right) Surface brightness profiles of sectors 1–10. Note that sector 7 contains NGC 4388, and sector 8 contains an X-ray clump locatednear NGC 4388.
Table 2.
Best-fit parameters of the β model fit for sectors 2, 3, 4, and 6. Parameter Unit Sector 2 Sector 3 Sector 4 Sector 6 β . ± .
03 0 . ± .
01 0 . ± .
01 0 . ± . θ (arcmin) . ± . . ± . . ± . . ± . S ( × − c s − arcmin − ) . ± .
06 1 . ± .
06 1 . ± .
10 1 . +0 . − . constant ( × − c s − arcmin − ) . ± . . ± . were introduced. In all the fitting, it was in the acceptable range( ∼ . ± . ). First, we analyzed spectra of outer regions to the south-eastand the south of M86 center, SE6 and S, respectively, to eval-uate the ICM around M86. One or two APEC models wereemployed to represent the emission in these regions, in addi-tion to two APEC models for the Galactic components (LHBand MWH) and a power law model for the CXB. When twoAPEC models were employed, their abundances were linked.The spectra and the best fit models are shown in figure 5, andthe best fit parameters are summarized in table 3. Fitting wassignificantly improved by employing two APEC models. The F -test provided the probability of . × − for region SE6and . × − for region S, respectively. In both cases, thehigher temperature component was dominant. The temperaturewas kT = 2 . +0 . − . keV and . ± . keV, respectively, andthe abundance was ∼ . Z ⊙ . The temperature of the othercomponent was kT ∼ keV. Since the higher temperature com-ponent was dominant, and its temperature was ∼ keV, it wasinterpreted as the ICM emission. The lower temperature com-ponent was, on the other hand, considered a contribution of theextended emission of M86. Note that the normalizations of theLHB and the MWH were significantly different between the two regions, and their error bars were large. We compared themwith those shown in Simionescu et al. (2015), who determinedthe spectral parameters of the LHB and the MWH using a set of12 ROSAT All-sky Survey data beyond the virial radius of theVirgo cluster. After the unit conversion, the normalization of theMWH of region S was consistent with Simionescu et al. (2015)within an error range, while that of region SE6 was larger by afactor of ∼ . The LHB normalizations were larger by about anorder even for region S. We also compared the normalizationswith those reported by Yoshino et al. (2009), who studied softX-ray diffuse foreground emission with Suzaku. The normal-izations of our results were within the variation range of thoseshown in Yoshino et al. (2009), except for the LHB normaliza-tion of region SE6. The foreground components of region SE6might have been affected by solar activities.A region around SE6 was observed with ROSAT and
Chandra . According to Rangarajan et al. (1995), the ICM tem-perature of the
ROSAT
South East quadrant was 1.76 keV, whileRandall et al. (2008) reported that the spectrum of their R18region was represented with kT = 1 . keV and 2.107 keVAPEC models. These were close to the temperatures of the T model and the T model shown in table 3, respectively.Therefore, we judged that the results were in agreement witheach other. On the other hand, the ICM temperature of ROSAT
South West quadrant was 2.09 keV (Rangarajan et al. 1995),which was not consistent with the temperatures of our region ublications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 (cid:0) . ( c oun t s s (cid:0) a r c m i n (cid:0) ) S u r f a c e b r i gh t ne ss Sector 20.1 1 10 . . r a t i o Angular distance from M86 center (arcmin) (cid:0) . ( c oun t s s (cid:0) a r c m i n (cid:0) ) S u r f a c e b r i gh t ne ss Sector 30.1 1 10 . . r a t i o Angular distance from M86 center (arcmin) (cid:0) . ( c oun t s s (cid:0) a r c m i n (cid:0) ) S u r f a c e b r i gh t ne ss Sector 40.1 1 10 . . r a t i o Angular distance from M86 center (arcmin) (cid:0) . ( c oun t s s (cid:0) a r c m i n (cid:0) ) S u r f a c e b r i gh t ne ss Sector 60.1 1 10 . . r a t i o Angular distance from M86 center (arcmin)
Fig. 3.
The best fit models for the surface brightness profiles of sectors 2, 3, 4, and 6.
Table 3.
Best fit parameters for the spectra of the outer regions SE6 and S.
Region SE6 Region SComponent Parameter Unit T model T model T model T modelAPEC1 kT (keV) 1.887 +0 . − . +0 . − . +0 . − . +0 . − . Z (solar) 0.199 +0 . − . +0 . − . +0 . − . +0 . − . Norm ( × − ) 4.191 +0 . − . +0 . − . +0 . − . +0 . − . APEC2 kT (keV) – 0.922 +0 . − . – 1.018 +0 . − . Z (solar) = APEC1 : Z Norm ( × − ) – 3.327 +5 . − . – 3.984 +1 . − . APEC (LHB) kT (keV) 0.11 (fixed)Norm ( × − ) 8.138 +2 . − . +2 . − . +0 . − . +0 . − . APEC (MWH) kT (keV) 0.3 (fixed)Norm ( × − ) 8.567 +5 . − . +6 . − . +2 . − . < . – χ / d.o.f. 333.37/279 320.99/277 494.53/424 472.16/422 The abundances of LHB and MWH were fixed at 1 solar. The normalizations of the APEC components are in units of − π [ DA (1+ z )]2 R n e n H dV per π arcmin , where D A is the angular diameter distance to the source (cm), n e and n H arethe electron and hydrogen densities (cm − ). Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 EX TailPlume Center S SE Fig. 4.
Regions for spectral analysis. Dotted circles show the distance fromthe M86 center, 5, 10, 15, 20, 25 arcmin, respectively.
S. In the following analysis, the T model parameters for re-gion SE6 was regarded as representative of the ICM emission,and the abundance of the ICM was assumed to be . Z ⊙ . Itis consistent with the number (26%) adopted by Randall et al.(2008). Note that the metallicity of the Virgo cluster at the sameradius from M87 was also reported to be Z ∼ . Z ⊙ (Urban etal. 2011; Ehlert et al. 2013). Secondly, the M86 center and the halo regions were analyzed.The regions used for the analysis were the center and 5 annularsectors (SE1–5, EX1–3) shown in figure 4. The center region isa circle of 1 . ′ θ shown in table 2, and surface brightness can be regardedapproximately constant in this radius. Annular sectors SE1–5were defined from the position angle 95 ◦ to 150 ◦ , to avoid con-tamination from NGC 4438 and NGC 4388. The radius of thefive regions were 1 . ′ . ′
5, 3 . ′ ′ , 6 . ′ . ′
5, 8 . ′ ′ , and 12 ′ –16 ′ , respectively. The dataset . ′ ′ , 5 ′ –10 ′ , and 10 ′ –16 ′ . Thedataset T ) modeland a two-temperature ( T ) model, represented by one or twovAPEC model(s), which is an APEC model with variable abun-dances, in addition to the background and foreground compo-nents described in the previous section. An additional power law was added to the center region, to represent the contribu-tion of unresolved low mass X-ray binaries. The photon indexwas fixed at 1.5 (e.g., Sarazin et al. 2003). To avoid the normal-izations of the galactic components varying region by region,we first fitted the center and SE1–6 regions simultaneously todetermine the normalizations of the galactic components, andfixed them at the values obtained in the simultaneous fitting.The fitting results of the T model and the T model are shownin table 4 for the center and SE1–5 regions, and in table 5 forEX1–3 regions. The spectra and the best-fit models of the T fitare shown in figure 6 for the center and SE1–5, and in figure 7for EX1–3. Since there was an uncertainty in the normalizationsof the LHB and the MWH as described in section 4.1, we inves-tigated how the results of the T fit were affected if these nor-malizations were fixed at the numbers obtained by Simionescuet al. (2015). When we fixed the normalization of the MWHat . × − , i.e., 1/10 of the number shown in table 4 and 5,the best-fit parameters were unchanged within a statistical errorrange even in SE4 and SE5. When we fixed the normalization ofthe LHB at . × − , i.e., 1.6% of the number shown in table4 and 5, the temperature of the vAPEC component of SE5 de-creased, the Fe abundances of SE4 and SE5 decreased, and thenormalizations of SE4 and SE5 increased, while other param-eters were unchanged within a statistical error range. In thesecases, however, reduced χ increased by 0.05 for SE4 and 0.11for SE5. When another APEC component was added, the tem-perature became close to that of the LHB. Therefore, a largernormalization of the LHB or equivalent was needed to repre-sent the Suzaku spectra. The temperatures of the T fit were about 0.8–1.0 keV. Asshown in the upper panel of figure 8, the temperatures of theinner regions (center and SE1) were lower, while they were al-most constant or slightly decreasing toward the outer regions(SE2–6). When the T fit and the T fit were compared, the F -test probabilities were . × − , . × − , . × − , . × − , . × − , and 0.37, for the center and the SE1–5 regions, respectively. Thus, the improvement of the fit wasreasonable except for SE5.In the center region, the temperature of the main compo-nent was . +0 . − . keV while that of the second componentwas ∼ . keV, when the T model was employed. The nor-malization of the second component was about 0.3 times thatof the first component. When the ICM temperature was madefree in the T fit, it resulted in ∼ . keV, rather than stayingaround 2 keV. Thus, the spectral data preferred the existence ofa cold component. Matsushita (2001) showed that the tempera-ture of the central region ( < . ′ ) was 0.69 keV, while Randallet al. (2008) reported existence of cold clumps around the core.Ehlert et al. (2013) pointed the presence of ∼ . –0.7 keV gas ublications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 (cid:0) . . C oun t s s e c (cid:0) k e V (cid:0) SE6 (2T) ✁ ✂ Energy (keV) (cid:0) . . C oun t s s e c (cid:0) k e V (cid:0) S (2T) ✁ ✂ Energy (keV)
Fig. 5.
Spectra of the outer regions SE6 and S, and the best-fit models. The red and black crosses show the data points of BI and FI CCD data, and thered and black solid curves are the best fit models for them. The green, magenta, and gray curves are the high- T component, the low- T component, and thebackgrounds/foreground components (CXB, LHB, MWH), respectively. Only the components for the BI model are shown. Table 4.
Best-fit spectral parameters for the center and SE regions obtained from the T fit and the T fit. Component Parameter Unit Center SE1 SE2 SE3 SE4 SE5 T modelvAPEC kT (keV) 0.800 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z O (solar) 0.666 +0 . − . +0 . − . +0 . − . +1 . − . +0 . − . < Z Ne (solar) 2.683 +0 . − . +0 . − . +1 . − . +2 . − . +1 . − . +2 . − . Z Mg (solar) 0.810 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . < Z Si (solar) 0.646 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z S (solar) 0.896 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z Fe (solar) 0.663 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Norm ( × − ) 16.288 +2 . − . +1 . − . +1 . − . +0 . − . +0 . − . +0 . − . APEC (ICM) kT (keV) 2.1 (fixed) Z (solar) 0.27 (fixed)Norm ( × − ) 2.720 +1 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Power law (LMXB) Γ × − ) 4.506 +1 . − . – – – – –– χ / d . o . f . T modelvAPEC1 kT (keV) 0.876 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z O (solar) 0.647 +0 . − . +0 . − . +0 . − . +1 . − . +0 . − . +0 . − . Z Ne (solar) 2.277 +0 . − . +0 . − . +1 . − . +3 . − . +0 . − . +2 . − . Z Mg (solar) 0.874 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z Si (solar) 0.715 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z S (solar) 0.921 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z Fe (solar) 0.745 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Norm ( × − ) 12.097 +2 . − . +1 . − . +0 . − . +0 . − . +1 . − . +0 . − . vAPEC2 kT (keV) 0.562 +0 . − . +0 . − . +0 . − . +0 . − . > . × − ) 3.625 +2 . − . +2 . − . +0 . − . +1 . − . +0 . − . +1 . − . APEC (ICM) kT (keV) 2.1 (fixed) Z (solar) 0.27 (fixed)Norm ( × − ) 3.048 +1 . − . +0 . − . +1 . − . +0 . − . +0 . − . Power law (LMXB) Γ × − ) 4.280 +1 . − . – – – – –– χ / d . o . f . In all the cases, APEC models for LHB ( kT = 0 . keV, Z = 1 Z ⊙ , Norm = 9 . × − ) and MWH ( kT = 0 . keV, Z = 1 Z ⊙ , Norm = 1 . × − ), and a powerlaw model for CXB ( Γ = 1 . , Norm = 1 . × − ) were included. The normalizations of the APEC components are in units of − π [ DA (1+ z )]2 R n e n H dV per π arcmin , where D A is the angular diameter distance to the source (cm), n e and n H are the electron and hydrogen number densities (cm − ). The normalizations ofthe power law are in units of photons keV − cm − s − at 1 keV per π arcmin . Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV) Center (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
SE1 (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
SE2 (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
SE3 (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
SE4 (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
SE5
Fig. 6.
Spectra of the center and SE1–6, and the best-fit models of the T fit. The red and black crosses show the data points of BI and FI CCD data, and thered and black solid curves are the best fit models for them. The blue, green, yellow, and gray curves are the vAPEC component, the ICM, the LMXBs, and thebackgrounds/foreground components (CXB, LHB, MWH), respectively. Only the components for the BI model are shown. ublications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 Table 5.
Best-fit spectral parameters for EX 1–3 regions obtained from the T fit and the T fit. Region EX1 Region EX2 Region EX3Component Parameter Unit T model T model T model T model T model T modelvAPEC1 kT (keV) 0.829 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z O (solar) 0.594 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z Ne (solar) 2.652 +0 . − . +0 . − . +1 . − . +0 . − . +0 . − . +0 . − . Z Mg (solar) 0.736 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z Si (solar) 0.531 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z S (solar) 0.589 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Z Fe (solar) 0.547 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . Norm ( × − ) 6.823 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . vAPEC2 kT (keV) – 0.606 +0 . − . – 1.461 +0 . − . – 0.630 +0 . − . Norm ( × − ) – 1.392 +0 . − . – 0.709 +0 . − . – 0.442 +1 . − . APEC (ICM) kT (keV) 2.1 (fixed) Z (solar) 0.27 (fixed)Norm ( × − ) 3.582 +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . +0 . − . – χ / d . o . f . In all the cases, APEC models for LHB ( kT = 0 . keV, Z = 1 Z ⊙ , Norm = 9 . × − ) and MWH ( kT = 0 . keV, Z = 1 Z ⊙ , Norm = 1 . × − ), and apower law model for CXB ( Γ = 1 . , Norm = 1 . × − ) were included. The normalizations of the APEC components are in units of − π [ DA (1+ z )]2 R n e n H dV per π arcmin , where D A is the angular diameter distance to the source (cm), n e and n H are the electron and hydrogen numberdensities (cm − ). The normalizations of the power law are in units of photons keV − cm − s − at 1 keV per π arcmin . (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
EX1 (1T) (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
EX2 (1T) (cid:0) C oun t s s e c ✁ k e V ✁ ✂ ✄ Energy (keV)
EX3 (1T)
Fig. 7.
Spectra of the center and EX1–3, and the best-fit models of the T fit. The red and black crosses show the data points of BI and FI CCD data, and the redand black solid curves are the best fit models for them. The blue, green, and gray curves are the vAPEC component, the ICM, and the backgrounds/foregroundcomponents (CXB, LHB, MWH), respectively. Only the components for the BI model are shown. Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 k T (cid:0) N o r m r (arcmin) Fig. 8.
Radial profiles of the temperature and the normalization of thevAPEC component of the T model of the center and SE1–6 regions. between M86 and NGC 4438. Our result was qualitatively con-sistent with them.For SE1, the temperature of the main component was almostunchanged while the second temperature was . +0 . − . keV.The normalization of the ICM, however, became 0, which wasunrealistic. If the ICM temperature was made free in the T fit, it became . +0 . − . keV, and the abundances increased by ∼ . Z ⊙ . The F -test probability was . × − , and hence,this improvement was reasonable. The results suggest that the T model is enough, but the ICM temperature could be as lowas ∼ . keV in SE1. Note that it is close to that of region Sshown in table 3.The second temperatures of SE2 and SE3 were 1.3 keVand 1.7 keV, respectively, suggesting the ICM temperature waslower than what we assumed like SE1. When the ICM tempera-ture was made free in the T fitting, they became . +0 . − . keVand . +0 . − . keV, respectively. In these cases, however, theabundances of the main component became unphysically large,and thus, the results were considered unacceptable. This isprobably because the room for the continuum became smaller,as the ICM temperature became lower. The results may sug-gest that the ICM temperature is located between them, but itwas not possible to further constrain them. For SE4, the secondtemperature was too high to constrain.As a conclusion, the T model is better for the center, whilethe T model is enough for SE1 but the ICM temperature couldbe as low as 1.6 keV. For SE2 and SE3, the ICM tempera-ture may also be lower than 2.1 keV, while that of SE4 maybe slightly higher.Note that the temperatures of EX1–3 showed a similar char-acteristics. When the T fit and T fit were compared, the F -test probabilities were × − , × − , and 0.023, forEX1, 2, 3, respectively. Thus, the improvement of the fit wasreasonable for EX1 and 2. The second temperature of EX1 was . +0 . − . keV, which may be the cooler component either lo-cated at the center or in the region between M86 and NGC 4438.The results of EX2 may indicate that the ICM temperature islower than what was assumed. The normalizations of the T fit were plotted as a function ofradius from the center, in the lower panel of figure 8. When theprofile was fitted with a β -model, the best-fit parameters were β = 0 . +0 . − . , θ = 2 . +2 . − . , and S = 0 . +0 . − . . β and θ were consistent with those of sector 4 shown in table 2 withinan error range.From the normalization of the APEC model, the emissionmeasure can be obtained. Assuming the center region as a uni-form sphere of 1 . ′ n H became ∼ . × − cm − , and the mass of thesphere M became ∼ . × M ⊙ , where mean molecularweight of hydrogen was assumed to be 1.4. When the nor-malization of the T model was adopted, the results were un-changed ( n H ∼ . × − cm − and M ∼ . × M ⊙ ). Notethat they are consistent with what was obtained by Randall et al.(2008) ( n core ≈ . × − cm − and M core ≈ . × M ⊙ within a sphere of radius 9.6 kpc). Figure 9 shows the radial distributions of the abundances of O,Ne, Mg, Si, S, and Fe of the center and EX1–3 regions. T fitresults of SE1–5 regions are also shown. The distributions wereconsistent with each other, although the errors of SE1–5 werelarge. Within r < ∼ ′ , the abundances were relatively large androughly constant, while outside r > ∼ ′ the abundances becamesignificantly smaller (0.2– . Z ⊙ ).Figure 10 shows the abundance ratios of elements with re-spect to Fe. O/Fe, Mg/Fe, Si/Fe, S/Fe are consistent with 1,while Ne/Fe is 2–4. When the T model was adopted, the abun-dances were generally slightly higher, by ∼ . Z ⊙ , but the over-abundance of Ne was unchanged. Thirdly, we fitted the spectra of the plume and the tail regionswith T model. The results are summarized in table 6 as case1 (ICM temperature fixed at 2.1 keV) and case 2 (ICM temper-ature free), and the best-fit models of case 1 are shown in fig-ure 11. The F -test probabilities between case 1 and case 2 were . × − and 0.20 for the plume and the tail, respectively. Theimprovement was reasonable for the plume, but the ICM tem-perature was > . keV in case 2, which seemed too high asthe ICM temperature. On the other hand, the abundances of thetail were high ( ∼ . Z ⊙ ) when the ICM temperature was fixedat 2.1 keV (case 1). The ICM temperature slightly higher than ublications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 . O ( s o l a r) Angular distance from M86 center (arcmin) 0.1 1 10 . N e ( s o l a r) Angular distance from M86 center (arcmin)0.1 1 10 . M g ( s o l a r) Angular distance from M86 center (arcmin) 0.1 1 10 . S i ( s o l a r) Angular distance from M86 center (arcmin)0.1 1 10 . S ( s o l a r) Angular distance from M86 center (arcmin) 0.1 1 10 . F e ( s o l a r) Angular distance from M86 center (arcmin)
Fig. 9.
Radial distributions of the abundances of O, Ne, Mg, Si, S, and Fe, as a function of angular distance from the M86 center. Black crosses and greensquares are T fit and T fit results of the center and EX1–3 regions, respectively. T fit results of SE1–5 are also shown with gray crosses. Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0
Table 6.
Best-fit parameters of the plume and the tail regions.
Plume TailComponent Parameter Unit Case 1 Case 2 Case 1 Case 2vAPEC kT (keV) 0.861 +0 . − . +0 . − . +0 . − . +0 . − . Z O (solar) 0.751 +0 . − . +0 . − . +0 . − . +0 . − . Z Ne (solar) 2.668 +0 . − . +0 . − . +2 . − . +1 . − . Z Mg (solar) 1.218 +0 . − . +0 . − . +0 . − . +0 . − . Z Si (solar) 0.813 +0 . − . +0 . − . +0 . − . +0 . − . Z S (solar) 0.990 +0 . − . +0 . − . +0 . − . +0 . − . Z Fe (solar) 0.906 +0 . − . +0 . − . +0 . − . +0 . − . Norm ( × − ) 10.230 +1 . − . +0 . − . +1 . − . +2 . − . APEC (ICM) kT (keV) 2.1 (fixed) > . +6 . − . Norm ( × − ) 2.095 +0 . − . +0 . − . +0 . − . +1 . − . – χ / d . o . f . In all the cases, APEC models for LHB ( kT = 0 . keV, Z = 1 Z ⊙ , Norm = 9 . × − ) and MWH ( kT = 0 . keV, Z = 1 Z ⊙ , Norm = 1 . × − ), and a power law model for CXB ( Γ = 1 . , Norm = 1 . × − ) were included. Thenormalizations of the APEC components are in units of − π [ DA (1+ z )]2 R n e n H dV per π arcmin , where D A is theangular diameter distance to the source (cm), n e and n H are the electron and hydrogen number densities (cm − ). Thenormalizations of the power law are in units of photons keV − cm − s − at 1 keV per π arcmin . . Z / Z F e ( s o l a r r a t i o ) Atomic NumberO/Fe Ne/Fe Mg/Fe Si/Fe S/Fe
Fig. 10.
Abundance ratios of O, Ne, Mg, Si, and S with respect to Fe.Triangles, circles, and squares represent the center, SE1–5, and EX1–3,respectively. T model, the temperature of the second compo-nent became too high to constrain. Therefore, T model wasnot meaningful for these regions.A similar abundance pattern to that of the center region wasseen in both case 1 and 2, i.e., the abundances of O, Mg, Si,S, Fe were close to each other, while that of Ne was about 2.5times larger.The abundance of Fe, especially in the tail region, wasslightly higher than that of the center. In the tail region, thesurface brightness was relatively low, and the abundance (alsothe normalization) was affected by the normalization and thetemperature of the ICM. Figure 12 shows a confidence contourbetween the Fe abundance and the ICM normalization when theICM temperature was fixed at 2.1 keV, and a confidence con- tour between the Fe abundance and the ICM temperature whenboth the ICM temperature and normalization were free. Thereis a positive correlation with the ICM normalization and a nega-tive correlation with the ICM temperature. The 90% lower limitof the Fe abundance is 0.9 if the ICM temperature is 2.44 keV.Therefore, we cannot conclude that the abundances in the tailregion is higher than those of the center region from the Suzakudata.If we assume a uniform prolate spheroid of the equatorialradius of 1 . ′ . ′ n plume ≈ . × − cm − and the total mass is M plume ≈ . × M ⊙ . This is consis-tent with the numbers obtained by Randall et al. (2008) withina factor of 2. If we assume a uniform cylinder of 1 . ′ . ′ n tail ≈ . × − cm − and the total mass is M tail ≈ . × M ⊙ . The mass is about 1/4 of that estimatedby Randall et al. (2008). Major difference is that our data onlycovered part of the tail. As shown in the previous section, the temperatures of the corewere kT = 0 . +0 . − . keV and ∼ . keV ( T fit), while those ofthe plume and the tail were . ± . keV and . ± . keV,respectively. Thus, the temperature of the tail was slightlyhigher. This is generally consistent with those reported byRandall et al. (2008) and Ehlert et al. (2013). There was a ten-dency that the abundances became slightly larger in the orderof the core, the plume, and the tail. However, we concludedthat the difference was not significant, thinking about statisti-cal errors and also about systematic errors due to variation of ublications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 (cid:0) C oun t s s e c ✁ k e V ✁ Plume ✂ ✄ Energy (keV) (cid:0) C oun t s s e c ✁ k e V ✁ Tail ✂ ✄ Energy (keV)
Fig. 11.
Spectra of the Plume and Tail regions, and the best-fit models. The red and black crosses show the data points of BI and FI CCD data, and the redand black solid curves are the best fit models for them. The blue, green, and gray curves are the vAPEC component, the ICM, and the backgrounds/foregroundcomponents (CXB, LHB, MWH), respectively. Only the components for the BI model are shown. . F e ICM Norm + . . . F e kT ICM (keV) + Fig. 12. (left) Confidence contour between Fe abundance and ICM normalization, when the ICM temperature was fixed at 2.1 keV. (right) That between Feabundance and ICM temperature, when both the ICM normalization and temperature were made free. The confidence levels are σ , 90%, and σ , respectively. the ICM temperature and normalization. Randall et al. (2008)pointed that the temperature structure of the tail is consistentwith a ram-pressure stripping model, i.e., the hotter, higher en-tropy group gas is stripped first, followed by the cooler, lowerentropy M86 ISM. Our results strongly support it, since the tem-perature of SE3 ( . +0 . − . keV) and the Fe abundance (0.7–0.8 Z ⊙ ) are close to those of the tail.We determined the abundances of O, Ne, Mg, Si, S, and Feseparately, and also showed that all the spectra of the differentregions had a very similar abundance pattern, i.e., O/Fe, Mg/Fe,Si/Fe, and S/Fe were basically consistent with the solar ratio,while Ne/Fe was larger by a factor of 3–4. Konami et al. (2014)analyzed Suzaku data of M86 as a whole, i.e., including thecore, the plume, and part of the tail together, and reported thatNe/Fe was about 3. Our analysis showed that it was the casewith the center, the plume, the tail, and also the halo. The spec-tra of these regions (figure 6 and 11) showed a peak at around 0.9 keV and a hump slightly above 1 keV. They were causedby a forest of Fe L lines, and Ne X Ly α lines at 1022 eV (Ne X p → s ), respectively, when the temperature was 0.8 keV. Thepeak energy due to the forest of Fe L lines rises as the temper-ature rises. Therefore, the spectral shape in this energy regionis mainly determined by the temperature and the abundancesof Fe and Ne. The Ne abundance by Lodders (2003) is about60% of that provided by Anders & Grevesse (1989) or Grevesse& Sauval (1998). In the recent solar abundances provided byLodders (2010), the Ne abundance is much closer to Anders& Grevesse (1989) or Grevesse & Sauval (1998). Even if thenew value were adopted, Ne overabundance would be still sig-nificant, by a factor of 2–3. This common abundance pattern,which is derived from the very similar spectral shapes, is oneevidence that the hot gas in the plume and the tail regions hasthe same origin as that in the core.Konami et al. (2014) reported that the large Ne/Fe ratio Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 cannot be explained by any mixture of SNe type Ia and core-collapse SNe, and concluded that Ne abundance may have in-trinsically large systematic errors because their emission linesare hidden by prominent Fe-L lines. We further investigatedit. To check model uncertainties, we fitted the center data withSPEX 3.0 (Kaastra et al. 1996). When the temperature wasfixed at the same number, the abundances differed by ∼ %.However, the difference of Ne/Fe was only about 5%, andhence, there was no significant difference between the resultsbased on APEC and SPEX. According to AtomDB, there is anFe L line (Fe XVII p d → p ) at 1023 eV. In addition, thereis a strong Fe L line (Fe XXI p d → p ) at 1009 eV, whoseemissivity is comparable to the summation of two Ne X Ly α lines when the temperature is ∼ . keV, and reaches the max-imum when the temperature is ∼ . keV. Therefore, it may bedifficult to determine the abundance of Ne precisely in this tem-perature range unless Ne X Ly α lines are separated from strongFe L lines. Note that Ji et al. (2009) showed high-resolutionspectra of M86 obtained with XMM-Newton reflection gratingspectrometers (RGS), and the Fe XXI p d → p line seemedto be resolved. However, the neon abundance was not reported. The extended emission around M86 was clearly detected withSuzaku. According to the XIS mosaic shown in figure 1, it ex-tends over 15 ′ (72 kpc) from the center in the east direction, andover 10 ′ (48 kpc) in the south-west direction. With the moder-ate spatial resolution of Suzaku, the surface brightness profileof the core and the extended halo was represented with a single β model of β ∼ . , as described in section 3, and it indicatedthat the emission spreads to ∼ ′ ( ∼ kpc). This picture isconsistent with the the spectral fit of SE6 region ( r = 16 ′ –20 ′ ),which showed the existence of a component of kT = 0 . keV.If the surface brightness extends to a certain distance follow-ing a β model, the actual extent of the gas must be significantlylarger. Therefore, our results strongly suggest that the halo ofM86 extends over 100 kpc, at least in the east direction.Using the effective radius r e of 1 . ′
74 (de Vaucouleurs etal. 1991), Suzaku detected X-ray emission up to ∼ . r e inSE1–6 regions. The ratio of the temperature at 4–8 r e to that at < r e was kT (4 − r e ) /kT ( < r e ) = 1 . +0 . − . , showing thepositive temperature gradient in the central region. Nagino &Matsushida (2009) denoted galaxies with the temperature ra-tio > . as X-ray extended galaxies and others as X-ray com-pact galaxies. According to their criteria, M86 is located inthe boundary area. The ratio of the stellar velocity disper-sion to the gas temperature β spec ≡ µm p σ /kT , where µ isthe mean molecular weight in terms of the proton mass m p , is β spec = 0 . , for the gas temperature of 0.9 keV and the stellar Fig. 13.
Gas mass and dynamical mass from the position angle 45 ◦ to 275 ◦ .Thin and thick curves correspond to the gas mass and the dynamical mass,respectively. velocity dispersion of 256 km s − (Roberts et al. 1991). Thisis close to the typical number of the X-ray extended ellipticalgalaxies (Matsushita 2001). These results suggest that the halogas is located in a larger scale potential structure than that ofthe galaxy, such as a galaxy group (Matsushita 2001; Nagino &Matsushida 2009). In section 3, we showed the surface brightness distribution isrepresented by a β model of β ∼ . . In this section, we estimatethe gas mass assuming the β model obtained in section 3. The β model was valid from the position angle ∼ ◦ to ∼ ◦ ,covering about 64% of the whole area, and hence, we used onlythis region.In section 4.2.2, we estimated that the hydrogen number den-sity in the core was (6 – × − cm − , assuming a uniformsphere of 1 . ′ r < . ′ )assuming the β model parameters shown in table 2, and de-rived the density at r = 0 . It became (4.0–6.3) × − cm − . Inthe following discussion, we assume that the hydrogen numberdensity at r = 0 is × − cm − .If the β model for the surface brightness distribution is validto infinity, the density is given by the following function: n ( r ) = n (cid:26) (cid:16) rr (cid:17) (cid:27) − β . (2)Since the hydrogen number density at the center is × − cm − , the density at r ∼ kpc is ∼ × − cm − .According to Urban et al. (2011), the electron density of theVirgo ICM is ∼ × − cm − at about 350 kpc from the cen-ter of the Virgo cluster. Hence, the densities of the halo gas andthe ICM are comparable at around r = 100 kpc.The mass of the halo gas was estimated by ublications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0 M gas = X i ∆ θ i π Z R πr ρ i ( r ) dr, (3)where R is the radius of the gas, ∆ θ i is the angle of the i -thsector, ρ i ( r ) is the mass density at r of the i -th sector. The gasmass thus obtained is shown in figure 13 as a function of radius.The mass of the halo gas in r < kpc became ∼ × M ⊙ .We further estimated the dynamical mass assuming the hy-drostatic equilibrium. The halo must be affected by the motionof M86 in the Virgo cluster, and also by the ram-pressure strip-ping, and hence, the hydrostatic equilibrium will not be a goodapproximation. However, we think it is still useful to discussthe condition of M86. We calculated the dynamical mass usingthe following equation: M dyn ( r ) = − kT rµm H G (cid:16) d ln ρd ln r + d ln Td ln r (cid:17) ≃ βkT rµm H G r r + r , (4)assuming the temperature is almost constant. The dynamicalmass thus obtained is also shown in figure 13. It became ∼ × M ⊙ for r < kpc. Then the ratio of the gas mass tothe dynamical mass M gas /M dyn became ∼ . .B¨ohringer et al. (1994) decomposed the X-ray surfacebrightness distribution obtained with ROSAT into several com-ponents, and estimated the total mass within 280 kpc is (1–3) × M ⊙ . If we use their number, M gas /M dyn becomesonly ∼ − . If we extend our calculation to 230 kpc (differ-ence of the distance corrected) assuming the same β models arevalid, the gas mass becomes ∼ × M ⊙ , and the ratio is ∼ − . Schindler et al. (1999) reported that the galaxy masswithin 240 kpc from M86 center is × M ⊙ , and the ratio ofthe galaxy mass to the total mass is 2–6%. Therefore, depend-ing on the actual spread of the gas, the ratio of the gas mass tothe galaxy mass also significantly differs, from ∼ . to ∼ .According to Sasaki et al. (2015), the gas mass fractions ofclusters to the hydrostatic mass are about 0.02–0.1. They alsofound that the ratio is − for several subhalos in the Comacluster, indicating significant fraction of the gas was removeddue to interaction with the ICM, such as ram-pressure stripping.Our results implicate that the gas mass to the dynamical massratio of M86 is − – − , suggesting it is also significantlyaffected by the interaction with the ICM. According to Forman et al. (1979), the ram-pressure strippingoccurs when the ram-pressure of the cluster gas exceeds theforce holding the gas in the galaxy: ρ ICM v > ρ ISM σ (5)where ρ ICM and ρ ISM are the ICM and ISM densities, and σ gal is the galaxy velocity dispersion (See also Gunn & Gott 1972).At the core of M86, the ram-pressure is ∼ eV cm − and ρ ISM σ is ∼ eV cm − , for n ICM = 2 × − cm − , n ISM =5 × − cm − , v = 1500 km s − , and σ gal = 256 km s − . Table 7.
Mean free path and travel time for closeCoulomb collisions.
Center HaloNumber density n H (cm − ) × − × − Mean free path λ ⊥ (kpc) 8.6 215Travel time (y) . × . × Therefore, the ram-pressure stripping condition is satisfied evenin the core.According to Frank et al. (1992), the mean free path λ ⊥ ofthe Coulomb collisions of two protons that causes the large de-flection ( ∼ ◦ ) is given by λ ⊥ ≈ m p v πe n , (6)and the time needed for the particle to travel the mean free pathis given by t ⊥ = λ ⊥ v ≈ m p v πe n . (7)They were calculated for the core and the halo and are shownin table 7. The mean free path is comparable to the diameter ofthe core and the halo, and hence, the close Coulomb collisionsoccur with a large probability. In fact, there are much moredistant scatterings, and the mean free path and the travel timewill be shorten by a factor of the Coulomb logarithm ( ∼ ).Time needed to strip all the halo gas is simply estimated bythe total number of the halo gas divided by the flux of the ICM,i.e., t strip ≈ π n ISM R πn ICM R v = 43 n ISM n ICM
Rv . (8)It becomes ∼ × y in the core, and ∼ × y in the outerhalo. Since the distance between M86 and M87 is 350 kpc inprojection, the crossing time is t cross > ∼ Rv ∼ × [y] . (9)Therefore, t strip < ∼ t cross , and hence, most of the gas in the coreand in the halo will be stripped if M86 passes through the Virgocluster center once.We showed that the gas of low metal abundance still remainsin the outer halo. On the other hand, it was indicated that the gasmass to the dynamical mass ratio is − – − , which suggestssignificant fraction of the halo gas has been stripped. The M86group is experiencing the stripping by the Virgo ICM right now. As shown in section 3, a faint X-ray clump was detected nearNGC 4388. The temperature of the gas was ∼ keV and itsflux in the 0.5–2 keV band was ∼ × − erg cm − s − .We could not find any literature mentioning it. We checkedthe NASA/IPAC Extragalactic Database (NED) to see if there http://ned.ipac.caltech.edu/. Publications of the Astronomical Society of Japan , (2017), Vol. 00, No. 0
VPC 0400VPC 0415VPC 0423SD SS J 122519.51+124618.0 W H L J 122512.2+142722SD SSCGB 658492X MM J 122514.3+123954CX O J 122521.2+124430CX O J 122529.5+124322
NGC 4388
IC 3303
Fig. 14.
XIS image around NGC 4388 and the X-ray clump. Circles, dia-monds and crosses correspond to a cluster or a group of galaxies, galaxieswith known redshift, and X-ray sources, respectively. are any associations with known background groups or clus-ters. One cluster WHL J122512.2 + z = 0 . andone group SDSSCGB 65849 were found near the XIS field ofview, but it is unlikely that either of them is an optical coun-terpart. Next, we looked for galaxies with known redshift andX-ray sources in NED. The results are summarized in figure 14.Among them, IC 3303 is located near the brightest part of the X-ray clump. The redshift of this galaxy is − . ± . ( Conselice, Gallagher & Wyse 2001), which is very close tothat of M86. Therefore, this clump might be part of the M86subgroup gas, though it is separated from the extended emis-sion around M86. Further study is needed. We analyzed the Suzaku data of M86 and its adjacent regionsto study the extended emission around it. The M86 core, theplume, and the tail extending toward the northwest were clearlydetected, as well as the extended halo around them. From theposition angle ∼ ◦ to ∼ ◦ , the surface brightness distri-bution was represented relatively well with a single β -model of β ∼ . up to 15 ′ –20 ′ .The Suzaku XIS spectra of the M86 center, the extendedhalo, the plume, and the tail were explained with one- ortwo-temperature plasma model, in addition to the Virgo ICMof kT ≈ . keV and other background/foreground compo-nents. The temperatures of the center were . +0 . − . keV and ∼ . keV. The temperatures of the core and the halo have a pos-itive gradient, and reach the maximum of kT ∼ . keV at r ∼ ′ or ∼ r e . Outside it, it is almost constant or slightly decreas-ing toward the outer regions. The temperature of the plume andthe tail were . ± . keV and . ± . keV, respectively.Therefore, the temperature of the tail is slightly higher than the core and the plume. These were qualitatively consistent withthe previous Chandra and XMM-Newton results (Randall et al.2008; Ehlert et al. 2013).We succeeded in determining the abundances of O, Ne, Mg,Si, S, and Fe separately, for the core, the plume, the tail, and thehalo for the first time. The best-fit values of the Fe abundance inthe core and in the plume were ∼ . , while that in the tail wasslightly higher ( ∼ . ). However, we cannot conclude that theabundance in the tail is higher, thinking about the normalizationand the temperature variation of the ICM. The abundance of thehalo is almost the same up to ∼ ′ , and then it becomes sig-nificantly smaller (0.2–0.3) at r > ∼ ′ . This means that the gasin the outer halo is less polluted by the metals produced in thegalaxy. In all the regions, the abundance ratios of O, Mg, Si,and S to Fe were ∼ , while Ne/Fe showed a significantly largernumber (2–4). This Ne overabundance is coming from the spec-tral features at around 1 keV, and is another evidence that theplume and the tail have the same origin as the core. However,the overabundance by a factor of 2–4 cannot be explained by theuncertainty of the abundances, or mixture of known SNe nucle-osynthesis models. Ne abundance may have intrinsically largesystematic errors as suggested by Konami et al. (2014).Our results suggest that the halo of M86 extends over100 kpc, at least in the east direction. The temperature at thecenter is slightly lower, and the ratio of the stellar velocity dis-persion to the gas temperature is only 0.47. These features in-dicate that the extended halo gas is located in a larger scale po-tential structure than that of the galaxy, such as a galaxy group(Nagino & Matsushida 2009; Matsushita 2001). Using the β models for sectors, we estimated the gas mass from the posi-tion angle ∼ ◦ to ∼ ◦ (64% of the whole area). It was ∼ × M ⊙ in r < kpc. If we further assume the hy-drostatic equilibrium, the dynamical mass in the same regionwas ∼ × M ⊙ , giving the ratio of the gas mass to the dy-namical mass M gas /M dyn ∼ . . If we adopt the dynamicalmass within 230 kpc provided by B¨ohringer et al. (1994), theratio becomes ∼ − . These ratios suggest the halo of M86is significantly affected by the interaction with the Virgo ICM.Simple estimation of the ram-pressure stripping lengthscalesand timescales showed that the mean free path is comparableto the size of the core or the halo, and the stripping timescale iscomparable or shorter than the crossing time through the Virgocenter. Therefore, most of the gas in the core and in the halowill be stripped if M86 passes through the Virgo center once.The fact that the low metal gas still remains in the outer halo in-dicates that the M86 group is experiencing the stripping by theVirgo ICM right now. Acknowledgments