Outflow in Overlooked Luminous Quasar: Subaru Observations of AKARI J1757+5907

0.86 MB·English

by  Kentaro Aoki

PASJ:Publ.Astron.Soc.Japan,1–??, vol. 63,Subaruspecialissue
(cid:13)c 2011.AstronomicalSocietyofJapan.
Outflow in Overlooked Luminous Quasar: Subaru Observations of
AKARI J1757+5907
∗
Kentaro Aoki,1 Shinki Oyabu,1,2 Jay P. Dunn,3 Nahum Arav,3 Doug Edmonds,3 Kirk T. Korista,4 Hideo
Matsuhara,5
and
Yoshiki Toba5
1Subaru Telescope, National Astronomical Observatory of Japan, 650 North A’ohoku Place, Hilo,
1
HI 96720, U.S.A.
1
0 2Graduate School of Science, Nagoya University, Chikusa-ku, Nagoya 464-8602
2 3Department of Physics, Virginia Tech, Blacksburg, VA 24061, U.S.A.
4Department of Physics, Western Michigan University, Kalamazoo, MI 49008-5252, U.S.A.
n
5Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency, 3-1-1 Yoshinodai, Sagamihara, Kanagawa
a
J 229-8510

3
2 (Received 2010 September10; accepted 2011 January 14)

Abstract
]
O
We present Subaru observations of the newly discovered luminous quasar AKARI J1757+5907,which
C showsanabsorptionoutflowinitsspectrum. Theabsorptionconsistsof9distincttroughs,andouranalysis
h. focuses on the troughs at ∼−1000 km s−1 for which we can measure accurate column densities of HeI*,
p FeII and MgII. We use photoionization models to constrain the ionization parameter, total hydrogen
- column density, and the number density of the outflowing gas. These constraints yield lower limits for the
o distance, mass flow rate and kinetic luminosity for the outflow of 3.7 kpc, 70 M yr−1, and 2.0 1043 ergs
⊙
tr s−1, respectively. Such mass flow rate value can contribute significantly to the metal enrichm×ent of the
s
intra-clustermedium. Wefindthatthismoderatevelocityoutflowissimilartothoserecentlydiscoveredin
a
[ massivepost-starburstgalaxies. Finally,wedescribethescientificpotentialoffutureobservationstargeting

this object.
1
Key words:galaxies: active—quasars: absorptionlines—quasars: emissionlines—quasars: individual
v
(AKARI-IRC-V1 J1757000+590759)
0
4
3
1. Introduction provements in analysis methods and in target selection
4
. (see discussionsinAravetal.2008;Dunnet al.2010),we
1
Roughly 20% of all quasars exhibit Broad Absorption published several more accurate determinations of these
0
Lines(BAL)intheirspectrum(Kniggeetal.2008;Hewett quantities (Moe et al. 2009; Dunn et al. 2010; Bautista
1
1 & Foltz 2003), which are indicative of outflows with ve- et al 2010; Arav et al 2010). Here we present a similar
: locities 103 104 km s−1. Kinetic energy and mass analysis of an outflow in a luminous overlookedquasar.
v ∼ −
emanatingfromquasarshavebecomekeyelementsinthe- We use the term “BAL outflow” to designate intrin-
i
X oretical modeling of the evolution of supermassive black sic absorption detected in the spectrum of a quasar (i.e.,
r holes (SMBH) and their host galaxies (e.g., Di Matteo et originatingfromoutflowingmaterialinthe vicinity ofthe
a al. 2005; Hopkins et al. 2008), the suppression of cool- AGN, see Barlow & Sargent 1997). The original BAL
ing flows in clusters (e.g., Ciotti et al. 2010; Bru¨ggen & definition (Weymann et al. 1991)was created to differen-
Scannapieco 2009), and enrichment of the intra-cluster tiate, in low-resolution spectra, AGN outflow absorption
and inter-galactic media with metals (e.g., Moll et al. systems from intervening absorbers that do not have a
2007). Collectively, harnessing quasars’ mechanical en- dynamical connection to the AGN and is now physically
ergy to help in driving the above processes is known as obsolete. Significant number of narrower absorption lines
“AGN feedback”. show intrinsic natures, i.e., time variability and partial
To assess the contribution of BAL outflows to AGN coverage of the background light source. The frequency
feedback scenarios, it is (at least) necessary to determine ofintrinsicabsorptionlinesarediscussedandsummarized
their mass flow rate (M˙ ) and kinetic luminosity (E˙ ). in Ganguly& Brotherton(2008). We thereforeuse “BAL
out k
Early attempts to do so were done by de Kool et al. outflow”todesignatethephysicalnature,ratherthanthe
(2001), de Kool et al. (2002), Hamann et al. (2001), and observational definition, of the phenomenon.
Wampler, Chugai, Petitjean (1995). Recently, using im- AKARI-IRC-V1 J1757000+590759 (hereafter
AKARI J1757+5907) was discovered during the follow-
∗ BasedondatacollectedatSubaruTelescope,whichisoperated up observations of AKARI mid-infrared (MIR) All-Sky
bytheNationalAstronomicalObservatoryofJapan.
Survey. The infrared satellite AKARI performed an
2 Aoki et al. [Vol. ,
all-sky survey at 9 and 18 µm as well as at four far- 2. Observations and Data Reduction
infrared bands (Murakami et al. 2007; Ishihara et al.
2.1. high resolution spectroscopy
2010). The initial identification of the AKARI MIR
All-Sky Survey sources involved association with the The high resolution spectroscopy of
Two Micron All Sky Survey (2MASS) catalog (Skrutskie AKARI J1757+5907 (R.A.=17:57:00.24,
et al. 2006). This search identified some AKARI MIR Decl.=+59:08:00.3 (J2000.0)) was done with the High
sources with F(9µm)/F(Ks) > 2 in the high galactic Dispersion Spectrograph (HDS; Noguchi et al. 2002)
latitude (b > 20◦) after excluding the sample in the attached to the Subaru 8.2 m telescope (Iye et al. 2004)
| |
LargeandSmallMagellanicClouds. AKARIJ1757+5907 on2010June 17(UT). The weather conditions werepoor
has a large ratio of mid-IR to near-IR flux density with thick cirrus and the seeing was unstable (>1.′′0).
(F(9µm)/F(Ks) = 11.1). This mid-IR source is also The slit width was set to be 1.′′0. The HDS setting
coincident with a bright near-UV and optical source. We was Yd covering between 4054 and 6696 ˚A. This results
present photometries of AKARI J1757+5907 in table 1. in a resolving power of R 36000. The binning was
∼
The photographic magnitudes are from USNO-B1.0 cat- 2 (spatial direction) 4 (dispersion direction). We
×
alog (Monet et al. 2003), and they are converted to g,r, obtained eight exposures of 1800 s each, however, one of
and i magnitude using equation 2 of Monet et al. (2003). them is unusable due to low signal-to-noise ratio.
We found the radio source NVSS J175659.82+590801.5 The data were reduced using IRAF1 for the standard
(6.5 0.5 mJy at 1.4 GHz) (Condon et al. 1998) is proceduresofoverscansubtraction,darksubtraction,cos-
±
coincident with AKARI J1757+5907. The ratio of radio mic ray removal and flat-fielding, where wavelength cal-
(5 GHz) to optical (4400 ˚A) flux density is 1.4 assuming ibration was performed using the Th-Ar lamp. the rms
radio spectral index is -0.5. The ratio indicates the wavelength calibration error is 0.011 - 0.013 ˚A. The one-
quasar is radio quiet (Kellermann et al. 1989). dimensional spectra were extracted from each exposure.
Thefollow-upspectroscopyusingKPNO2.1mtelescope Heliocentric correction was applied. After that, all the
revealed that AKARI J1757+5907 is a z =0.615 quasar spectral exposures were combined, and all orders were
that shows HeI* and MgII absorptionlines as well as Hβ connected to one spectrum. We normalized the spectrum
and strong FeII emission lines (Toba et al, in prepara- using spline fits. Finally we converted to vacuum wave-
tion). The spectrum resembles the one of QSO 2359- length scale.
1241 (Brotherton et al. 2001; Arav et al. 2001). Both
2.2. Spectrophotometry
quasars show rare HeI* absorption as well as MgII ab-
sorption. They also have redder continua and strong The low resolution spectrophotometry of
FeII emission lines. The high resolution spectroscopy of AKARIJ1757+5907wasdoneon2010June30(UT)with
QSO 2359-1241 by Arav et al. (2001) revealed FeII ab- FOCAS (Kashikawa et al. 2002) attached to the Subaru
sorption lines, Thus FeII absorption lines are expected telescope. The new fully-depleted-type CCDs developed
to exist in AKARI J1757+5907. Following Korista et al. by NAOJ/ATC and fabricated by Hamamatsu Photonics
(2008),by measuring the ionic column densities (N ) of K. K. were installed and commissioned at that time. We
ion
FeII and HeI*, we can constrain the ionization parame- obtainedsixspectraof5minutesintegrationunderaclear
ter (U ) and hydrogen column density (N ). The high condition and good seeing (0.′′6 - 0.′′9 ). The 2.′′0 width
H H
brightness of this quasar permits us to do high resolution slit was used for spectrophotometry purpose. We used
spectroscopy of HeI* absorption lines and searchfor FeII two configurations: the R300 grism with the O58 filter
resonance and excited state absorption lines. The N (‘red’) and the B300 grism without any order-cut filters
ion
ratio of FeII* to FeII(E=0) yields the hydrogen number (‘blue’). The first three 300sspectra were the ‘red’ which
density, which in turn yields the distance of outflowing covers between 5700 ˚A and 10200 ˚A. The last three 300s
gas from the central source, E˙ and M˙ by using the spectra were the ‘blue’ which have an uncontaminated
k out
number density, U and N . range between 3500 ˚A and 7000 ˚A. The atmospheric
H H
The plan of this paper is as follows. In 2 we describe dispersion corrector was used. The slit position angle
the observations and data reduction. In§ 3 we deter- was 0◦, and the binning was 2 (spatial direction) 1
§ ×
mine the redshift of the object. The outflow absorption (dispersion direction) The spectrophotometric standard
troughs are discussed in 4 and the spectral energy dis- starBD+28◦4211wasobservedforsensitivitycalibration.
§
tribution in 5. In 6 we describe our photoionization The data were reduced using IRAF for the stan-
§ §
modeling, and in 7 the resultant determination of mass dard procedures of bias subtraction, wavelength cali-
§
flow rate and kinetic luminosity for the outflow. In 8 bration, and sky subtraction, except for flat-fielding.
§
we discuss our results and in 9 we describe the scien- AKARI J1757+5907 is so bright that its counts on
§
tific potential of additional Subaru/HDS and HST/COS the CCD are comparable to the flat frames, and are
observations targeting this object. Throughout this pa- much higher at shorter wavelength (<4000 ˚A). The flat-
per we assume H =70 km s−1 Mpc−1, Ω =0.3, and fielding procedure significantly reduced its signal-to-noise
0 m
Ω =0.7. Note that wavelengths of any transition in this
Λ 1 IRAF is distributed by the National Optical Astronomy
paper are ones in vacuum, and the observed wavelength
Observatory, which is operated by the Association of
scale is converted to one in vacuum.
UniversitiesforResearchinAstronomy,Inc.,undercooperative
agreementwiththeNationalScienceFoundation.
No. ] AKARI J1757+5907 3
ratio. We therefore skipped the flat-fielding procedure. 0.6155. We thus adopt 0.61525 0.00025 as the systemic
±
WavelengthcalibrationwasperformedusingOHnightsky redshift of the quasar. The blue component of [OIII] is
emissionlines for the redspectra andthe Cu-Ar lamp for blueshifted by 980 km s−1 relative to the systemic red-
the blue ones. The rms error of wavelength calibration is shift.
0.2 ˚A. The seeing was much smaller than the slit width,
thustheresolutionwasdeterminedbytheseeingdisksize. 4. Absorption lines
The resulting HeI* 3889 absorption line is 13 ˚A width at
6260 ˚A. This value corresponds to the resolving power of In the HDS spectrum of AKARI J1757+5907, the ab-
480( 620kms−1), whichis similar to the resolutionob- sorption lines are not heavily blended. The identifica-
tained∼with the 0.′′8 width slit, and is consistent with the tion is thus a straightforward task. We identify MgII
seeing size during our observations. The sensitivity cali- λλ2976,2803, HeI* λ2945, λ3188, λ3889, FeII λ2600,
bration was performed as a function of wavelength. The λ2586 as well as weaker absorption troughs from MgI
flux of the blue and red spectra at the same wavelength λ2852, HeI* λ2829, and CaII λλ3934,3969. The strong
agree within 1.6%. The foreground Galactic extinction absorptionlines suchastheMgIIdoubletandHeI*λ3889
of E(B V)=0.043 mag (Schlegel, Finkbeiner, & Davis clearly show 9 distinct troughs, which span a velocity
1998) w−as corrected. range from 660to 1520km s−1. The troughat 1000
km s−1 has−the sam−e outflow velocity as the blue−com-
3. Results of spectrophotometry ponent of the [OIII] emission line. We show the ab-
sorption troughs from MgII λλ2976,2803, HeI* λ2945,
Figure 1 displays the low-resolution optical spectrum λ3188, λ3889, and FeII λ2600, λ2586 in figure 3. We
of AKARI J1757+5907. The spectrum shows strong ab- do not detect an FeII λ2612 absorption trough from the
sorptionlinesofMgIIandHeI*aswellasemissionlinesof E=385cm−1 excitedlevel,whichhasthelargestoscillator
MgII,FeII,Hγ,Hβ,[OIII]. Inordertodeterminethesys- strength of the FeII* lines from this energy level, in our
temic redshift of the quasar, we deblend the [OIII] emis- spectralrange. WeshowthespectralregionofFeIIλ2612
sion lines from the Hβ and FeII emission lines. As seen in the lower panel of figure 3.
in figure 1, AKARI J1757+5907has strong FeII emission
4.1. Column Density Determinations
lines. However, the FeII emission template around Hβ
derived from the spectrum of I Zw 1 (Aoki, Kawaguchi, Both the MgII and HeI* troughs span the full veloc-
& Ohta 2005) is not a good match to the FeII emission ity range from 660 to 1520 km s−1 and appear in all
− −
in AKARI J1757+5907 (figure 2a). The intensity ratios 9 distinct troughs (see Figure 3). There is self blend-
among FeII emission lines are different between these ob- ing in the MgII troughs at the velocity extremes, which
jects. Thus, we cannot use the FeII emission template does not affect the majority of the troughs. The HeI*
derived form the spectrum of I Zw 1. and FeII troughs are free of any self blending. Of the 9
Insteadwefitthespectrumbetween7662˚Aand8204˚A troughs,onlythreehavecorrespondingabsorptioninFeII.
with a combination of FeII λλ4925,5020, Hβ, and [OIII] These are in the range of 1050 to 800 km s−1 (see
− −
after subtraction of a power-law continuum. This power- Figure3). Thus,weconcentrateonthisvelocityrangefor
law continuum is constructed by fitting at 6526-6544 ˚A ourmeasurementsasthebestphotoionizationconstraints
and 8842-8891 ˚A. The Hβ emission line is fitted with a are achieved by contrasting HeI* and FeII column densi-
combination of three Gaussians. We fit the [OIII] dou- ties (see section 6; Korista et al. 2008; Arav et al. 2010).
bletλλ4960.3,5008.2withtwosetsoftwoGaussians. The To determine the ionic column densities, we rebin the
widthandredshiftareassumedtobethesameforeachset, data to 10 km s−1 and use the velocity dependent ap-
andtheintensityratioisfixedtobe3.0. FeIIλλ4925,5020 parent optical depth (AOD), covering factor (C(v)), and
aremodeled by twoLorentzianswiththe same width and power-law fitting methods from Dunn et al. (2010; meth-
redshift. Their redshift is fixed to be the same as the ods1,2,and3 intheir subsection3.2)acrossthe rangeof
strongest Gaussian component of the Hβ. The result of 1050 to 800 km s−1.
− −
fitting is shown in figure 2b. We need a separate “red The AOD methods assumes that the emissionsource is
wing” of Hβ to get a satisfactory fit. This component completely and homogeneously covered by the absorber,
may be [OIII] emission line. The derived redshifts and so that the optical depth (τ(v)) at a given velocity is re-
FWHMs are tabulated in table 2. The FWHM is cor- lated to the normalized intensity via: I(v)=exp( τ(v)).
−
rectedfortheinstrumentalbroadeningbyusingthesimple The covering factor method assumes that at a given ve-
assumption: FWHM = (FWHM2 FWHM2 )1/2, locity, a fraction C(v) of the emission source is covered
true obs− inst
where FWHM is the observed FWHM of the line and with a constant value of optical depth, while the rest
obs
FWHM is aninstrumentalFWHM (620 kms−1). The of the source is uncovered. In order to for both τ(v)
inst
redshift ofthe red componentof [OIII] is 0.6150 0.0001. and C(v) we use use at least two absorption lines from
±
Therestequivalentwidthof[OIII]includingbothcompo- the same energy level of the same ion, and solve for
nentsis9.9 0.3˚A.Thisvalueisconsistentwiththe[OIII] I(v)j =1 C(v)+C(v)exp( τ(v)j), where Ij is the nor-
± − −
strength in majority (>50 %) of MgII BAL quasars re- malized intensity of the absorption due to the j tran-
ported by Zhang et al. (2010). We also detect weak [OII] sition in the same energy level. The ratio of different
emission at 6023.4 ˚A, which corresponds to a redshift of τ(v) are known from atomic physics and therefor the
j
4 Aoki et al. [Vol. ,
set of equations is solvable. The power-law model as- intensity at MgII λ2796 and λ2803 trough, respectively.
sumes that the absorption gas inhomogeneously covers Thus, the real residual intensity may be smaller than the
the background source. The optical depth is described observed one. Therefore, we conclude that the MgII col-
by τ (x)=τ (v)xa, where x is the spatialdimensionin umn density can be much larger than the AOD or C(v)
v max
the plane ofthe sky, a isthe powerlaw distributionindex determined values.
and τmax is the highest value of τ at a given velocity. In Finally,weshowtheresultforFeIItroughsinFigure6.
caseofthepowerlawmodel,τ isaveragedoverthespatial Unlike MgII and HeI*, the FeII troughs are both shallow
dimension x. and in a much lower signal-to-noise region of the spec-
In order to convert τ in each model to column density trum (towards the short wavelengthend of the detector).
we use Due to this, we find that the weaker FeII λ2587 line is
1 only detected across three velocity bins. Using both the
N(v)=3.8×1014fλτ(v)(cm−2km−1s), FeII λ2600 and the FeII λ2587 lines we calculate NFeII
for both C(v) and power-law methods for the three bins
whereλisthewavelengthofthelinein˚A,f istheoscilla- and include them in the integrated total. Due to a lack
torstrength. We use the oscillatorstrengthsfromFuhr & of detection of the λ2587 line, we use the λ2600 line to
Wiese (2006)for the FeII lines and values from the NIST calculate the column density from the AOD method for
Atomic Spectra Database (2010) for MgII and HeI*. theremainingpoints. Wealsocheckedthecolumndensity
InFigure4,weshowtheHeI*troughsandcolumnden- for FeII using the covering factor of MgII. The column
sity determinationfrom the three lines of HeI* presentin density calculated in this fashion changed only by 10%
∼
the HDS spectrum. The HeI* lines are well separated, as MgII nearly fully covers the source (C(v) 0.9) in this
≈
have no self blending, and we obtain consistent results velocity range.
with all three column density extraction methods. There
4.2. Column Density Limits on FeII* Excited State
aretwo velocitypoints wherethe C(v) becomes nonphys-
ical (i.e., negative or larger than 1.0), near 930 km s−1 Lines
and at velocities lower than 850 km s−1.−This occurs In order to help constrain the photoionization models
−
becausethe twoweakerlinesareveryshallowattheseve- insection6,weestimatethecolumndensitylimitsforthe
locities and are thus consistent with the continuum level FeII* λ2612andλ2757lines fromthe 385and7955cm−1
and dominated by the noise. energylevels,respectively. Neitherlineshowsadetectable
We show in Figure 5 the trough profiles and column troughinthe data. Therefore,wecanuse thetroughpro-
density results for MgII. Both the red and the blue dou- file of FeII λ2600 to determine the relative optical depth
blet troughs are quite deep and therefore their level of andcolumndensityofthesetwoenergylevels(seeSection
saturation is model dependent. The C(v) solution sug- 3.3ofDunnetal.2010). Wefindupperlimitsontheionic
gests a small level of saturation (only 40% larger column column densities of (3.7 and 0.5) 1012 cm−2 for the 385
density than the AOD estimate), while the power-law re- and 7955 cm−1 energy levels, resp×ectively.
sult is three times higher. This is an inherent feature
of the absorption models where in order to fit deep dou- 5. Determination of the Spectral Energy
blettroughsthepower-lawmodelrequiresamuchgreater Distribution
column density than the C(v) solution (see the case of
the OVI doublet in the spectrum of Mrk 279, Arav et al. Our FOCAS spectrophotometry data and the GALEX
2005). Since we do not have data for troughs from ad- photometry clearly show the flux drops at shorter wave-
ditional MgII lines, we cannot determine which model is lengths (<2000˚A inthe restframe)andindicate redden-
more physical. ing by dust (figure 7). We must consider extinction for
There is a possibility that a narrow MgII emission line deriving the intrinsic spectral energy distribution (SED).
fills the trough. As already noted early in this section, First, we measured the continuum flux at four points
the blue componentof[OIII]has the velocity of-1000km where there are less FeII emission contaminates. These
s−1. The MgII emission line from the same gas probably four points and the GALEX photometry points are then
exists. We estimate the flux of MgII emission line using shifted to the rest frame. We fit a reddened power-law
thefluxratioof[OIII]λ5008toMgIIinSeyfert2galaxies, to the continuum points. We adopted the index of the
NGC 1068 (Kraemer et al. 1998) and Mrk 3 (Collins et power-law continuum (α; f να) of the LBQS compos-
ν
∝
al. 2005). The ratios vary along the physical position in ite, -0.36 measured by Vanden Berk et al. (2001). We
thegalaxiesbetween0.03and0.15,andaverageofextinc- adopt the SMC-type extinction law. The best fit gives
tion corrected ratios are 0.09, and 0.07 in NGC 1068 and us the color excess E(B V) of 0.18. The best fit of the
−
Mrk 3, respectively. We also assume a 1:2 ratio for the reddened power-law continuum is shown in figure 7.
MgII doublet, and gaussians of the same width (σ=205 Toestimatethedistancetotheoutflowfromthecentral
km s−1) as the [OIII] emission line. This width corre- source (R), we need to determine the flux of hydrogen
sponds to 3.0 ˚A at the MgII observedwavelengthof 4500 ionizingphotonsthatirradiatestheabsorber(seeequation
˚A. The HDS spectrum before normalization is scaled to [1] below). Using the Mathews & Ferland (1987) SED,
the low-resolution spectrum. The expected height of the reddened to match the observed spectrum, we find that
MgIIemissionlinewillbe 13%and 4%oftheresidual number of hydrogen ionizing photons emitted per second
∼ ∼
No. ] AKARI J1757+5907 5
by the reddened centralsource (Q ) is 2.2 1057 photons Since the data do not provide a lower limit to the elec-
H
s−1. Here we assumed that the reddening×occur between tron number density, we find other valid solutions by re-
the central source and the outflow, as is the case where ducing n . When the hydrogen number density is re-
H
the photons are attenuated by the edge of the putative duced, the HeI* population drops. This is due to the
AGN obscuring torus (see full discussion in Dunn et al. factthatHeI*ispopulatedbyrecombinationofHeIIand
2010). This assumption will also give us smaller values the number of recombinations per unit time depends lin-
for the inferred R and therefore conservative lower limits early on n . Therefore, N must increase in order to
e H
for M˙out and E˙k. The LBol for the dereddened, intrinsic provide enough HeI* to be consistent with the measured
spectrum is 3.7 1047 erg s−1. value. FeII becomes dominant near the hydrogen ioniza-
×
tion front, while HeI* drops off drastically at the front.
6. Photoionization Modeling Thus, solutions in this region of the slab have U fixed
H
by N and the ratio N /U fixed by N . Due
He I∗ H H Fe II
Through photoionization modeling, reliable measure- to the tight correlationof N and U , all valid solutions
H H
mentsofHeI*andFeIIcolumndensitiesprovideaccurate have a similar ratio laying (nearly) on a straight line in
constraints on the total hydrogen column density, N , theN -U planefornumberdensitiesgreaterthan 100
H H H
and the hydrogen ionization parameter, cm−3. At these densities, the slabs do not form a h∼ydro-
gen ionization front. As we go to densities lower than
Q
UH ≡ 4πR2HnHc, (1) ∼inc1r0e0ascemli−n3e,aralyfrwonitthfoNrm.s, behind which FeII and MgII
H
where R is the distance from the central source, nH is In order to determine what effects choice of SED may
the total hydrogen number density, and c is the speed haveonN andU ,wecompareresultsoftheMF87SED
H H
of light. We use version c08.00 of the spectral synthesis with results of a softer SED. The soft SED has an opti-
code Cloudy, last described by Ferland et al. (1998), to cal to X-ray spectral index α = 1.5 compared to the
ox
modelaplane-parallelslabofgaswithconstanthydrogen MF87SEDwithα = 1.4(withth−econventionF =να)
ox ν
number density irradiated by a source continuum. We andwasgeneratedusin−gtheCloudycommandagn375000
focus on the kinematic components spanning a velocity -1.50 -0.125 -1.00, where the numbers are the tempera-
range from 800 km s−1 to 1050 km s−1 where we de- ture of the UV bump, α , α , α , respectively. We
− − ox uv x
tect FeII(0). In this velocity range, measurements of up- find the resulting N and U are nearly identical and
H H
perlimitsontheFeII(E=385cm−1)andtheFeII(E=7955 conclude that changing the SED only affects the energet-
cm−1) yield upper limits of electron number density ne, ics through QH, a finding consistent with the analysis of
103.6 and 103.8 cm−3, respectively. We adopt the conser- QSO 1044+3656reported in Arav et al. (2010).
vative value of an upper limit of ne 103.8 cm−3, and Anotherassumptioninourmodelsis solarabundances.
≤
assume nH ne, which is valid within the ionized zone To check the sensitivity of our results to metallicity
≈
we are discussing. While the electron number density is changes, we use the abundances in table 2 of Ballero et
well constrained from above, there are no diagnostics for al.(2008)formetallicityZ=4.23withthe MF87SEDand
a lower limit on ne in the data. lognH =3.8. While helium abundances are expected to
We begininvestigationofthe parameterspaceby using increase with oxygen abundances (e.g., Olive & Scully
Cloudy’s optimization mode to determine NH and UH (1996)), the amount of increase varies for different galax-
for nH =103.8 cm−3. The parameters NH and UH are ies. Therefore, to be conservative,we increase the helium
varied and ionic column densities are computed for each abundance significantly to 15% abovesolar. We find that
set of parameters. Best fit values are determined by χ2 logU is approximately 0.02 dex lower and logN is ap-
H H
minimization for given tolerances in the measured ionic proximately 0.15 dex lower for the increased metallicity
column densities. We adopt the measured ionic column modelwithlog(N /U )=22.83. We discussthe effectof
H H
densities determined by the partial coveringmethod, and these changes on the energetics of the outflow in the next
optimize to NFe II and NHe I∗, since these are the more section.
robustmeasurements. Themeasuredandmodelpredicted
columndensitiesarepresentedintable3. Forsolarabun- 7. Energetics of the Outflow
dances and the MF87 SED, as implemented by Cloudy2,
we find log UH = 2.15 and log NH =20.82 yield good Of particular interest for any outflow are its mass
−
fitstothecolumndensitiesofHeI*andFeII whileNMg II (Mout), the average mass flow rate (M˙out) and mechani-
is overpredicted. However, as discussed in section 4, the calworkoutputorkineticluminosity(E˙ ). Assumingthe
k
MgIItroughsmaybemoresaturatedthanthepartialcov- outflow is in a form of a thin partial shell moving with
ering model suggests. A hydrogenionization front, which
a constant radial velocity (v), at a distance R from the
we define as the positionat whichhalf of the totalhydro-
source, the mass of the outflow is:
gen is neutral (approximatedby N =1023U ), does not
H H
forminthis solutionalthoughwe areveryclose to itwith Mout=4πµmpΩR2NH, (2)
log(N /U )=22.96.
H H where N is the total column density of hydrogen, m is
H p
2 ThisSEDdiffersfromtheMF87SEDbytheadditionofasub- the mass of a proton, µ=1.4 is the plasma’s mean molec-
millimeterbreakat10µm. ular weight per proton, and Ω is the fraction of the shell
6 Aoki et al. [Vol. ,
occupied by the outflow. The average mass flow rate is and the relativistic effects. The M˙ is 110 M yr−1,
acc ⊙
given by dividing the outflowing mass by the dynamical which is similar to the lower limit of M˙ .
out
time scale of the outflow R/v (see full discussion in Arav WenotethattheΩ=0.2weuseisbasedthepercentage
et al 2010), therefore of quasars showing CIV BALs. In a recent work, Dai et
al. (2010) have shown that in the near infra-red surveys
M
M˙out out =4πµmpΩRNHv (3) low ionization BALs (LoBALs) fraction is 4%, consider-
∼ R/v
ably higher than deduced from optical surveys (probably
and E˙ = 1M˙ v2 2πµm ΩRN v3. (4) due to obscuration effects in the optical band). A more
k 2 out ∼ p H appropriate comparison in our case is to include some-
For the troughs we consider in AKARI J1757+5907, what narrower outflows with 1000 km s−1 <∆v <2000
the median velocity of the system is 970 km s−1. We km s−1. For these LoBALs based on “Absorption index”
assume that Ω = 0.2, which is the per−centage of quasars (Trumpetal.2006),Daietal.(2010)finda7.2%fraction.
showing BALs in their spectrum (see discussion in Dunn The frequency of HeI* outflows is much less known. The
et al. 2010). strongest HeI* line in the optical (λ3889) is shifted out-
To determine the distance, we solveequation(1) for R, sidetheopticalrangeforobjectswherewecandetectCIV
whichdependsonUH andnH. ThelackofFeII*detection λ1550 from the ground (e.g., z=1.5 for SDSS spectra), so
yields an upper limit of n <103.8 cm−3 (see Section 6), a meaningful census of these outflows is difficult to come
H
and therefore, as shown below, a lower limit on R. by. Anecdotally, we find that in most cases where we de-
For the range 101.8<n <103.8, our photoionization tect FeII absorption trough, we also detect HeI* troughs
H
solutionsobeythe relationshipsN U andU n−α, provided we have a clear spectral coverage of the latter
where0.4<α<1(αdecreasesasnH∝incHreases).HT∝hefiHrst (e.g., Arav et al. 2008; Arav et al. 2010).
H
relationshiparisesfromtherequirementofbeingclosetoa We also point out that AKARI J1757+5907 as well as
hydrogenionizationfront,andthesecondisduetothede- QSO 2359-1241 do not have measurable FeII absorption
creasingelectronpopulationattheHeI*metastablelevel in their low resolutionspectra. Also, their outflow veloci-
forlowernumberdensities. Therefore,equation(1)yields ties( 1200kms−1)andwidths oftroughsaremoderate.
∼
(α−1)/2 Quasars with similar moderate outflows are more numer-
R n . We use the solar abundances photoion-
∝ H ousthanextremeFeLoBALs(e.g.,SDSSJ0318-0600,Hall
ization solution values from Section 6 (logU = 2.15,
logN = 20.82 cm−2 for n =103.8 cm−3) anHd Q−from etal.(2002)),andasweshowhere,canhavesimilarmass
H H H flowrateandkineticluminosityasthemoreextremeones.
Section 5. Inserting these values into equation (1), pro-
Thisfactsuggeststhemassflowrateandkineticluminos-
vides a lower limit on the distance, R > 3.7 kpc for
ity values found here are more common among quasars,
n = 103.8 cm−3, which only increases to 6.6 kpc for
H than judged by the rarity of extreme FeLoBALs. If we
n =101.8(cm−3) due to the weak dependence of R on
H assume as a conservative limit that the Ω of the outflow
n . From equation (3) we observe that M˙ and E˙ de-
H out k seeninAKARIJ1757+5907issimilartothe7.2%LoBALs
pend linearly onthe product RN , whichis proportional
H fractionfoundbyDaietal.(2010)thenΩ =0.36,which
to n−(1+α)/2. Therefore, the upper limit for n provides 0.2
H H willreducethe valuesfor massflowrateandkinetic lumi-
lowerlimit to M˙out andE˙k. Using the NH andR derived nosity accordingly.
for n =103.8 cm−3, we find lower limits of M˙ > 70
H out
Ω0.2 M⊙ yr−1 and E˙k > 2.0 1043 Ω0.2 ergs s−1, where 8. Discussion
×
Ω Ω/0.2.
0.2
In≡theprevioussection,wediscussedchangesinNH and In Table 4 we showour M˙out andE˙k determinations in
UH duetometallicityandSEDchanges. UsingtheBallero quasarBALoutflowstodate(theolder,lessaccurate,ones
etal.(2008)abundancesfor Z/Z⊙=4.23reducesthe mass areshownintable10ofDunnetal2010). Whilethelower
flow rate by 30%. Changing to the soft SED described limit on E˙ for the AKARI J1757+5907outflow is rather
∼ k
intheprevioussectionresultsinverysmallchangesinNH low for AGN feedback purposes, we note that the corre-
and UH, but QH increases by a factor of ∼2, increasing sponding M˙out value is large enough to yield a significant
the mass flow rate by a factor of √2. contribution for the metal enrichment of the intra-cluster
∼
The mass of the black hole (MBH) of mediumaroundtheparentgalaxy(seeHallmanandArav
AKARI J1757+5907 is derived to be 4.0 109M⊙ 2010 [ApJ submitted])
×
based on the width of the Hβ emission line of σ =2160 We also note that this moderate velocity outflow is
kms−1 andthedereddenedopticalluminosityλL(5100˚A) similar to those discovered in post-starburst galaxies at
of 3.8 1046 erg s−1. We use the formula in Bennert et z 0.6 (Tremonti, Moustakas, Diamond-Stanic 2007).
al. (20×10) based on the calibrations of broad-line region Th∼ose galaxies show an outflow with MgII absorption
size-luminosity relation (Bentz et al. 2006) and the virial and a velocity of 1000 km s−1 . Their stellar masses are
coefficient taken from Onken et al. (2004). The derived as high as (0.7 4.8) 1011M (calculated from stellar
⊙
logLbol/LEdd = 0.13 is used for calculation of mass population synt−hesis m×odeling fit to their spectra). For
−
accretion rate (M˙acc) based on the accretion disk model comparison, the MBH of AKARI J1757+5907 is 4.0
byKawaguchi(2003),whichtakesintoaccounttheeffects 109M , therefore based on M -M relation (H¨arin×g
⊙ BH bulge
of electron scattering (opacity and disk Comptonization) & Rix 2004), the bulge of its host galaxy is estimated
No. ] AKARI J1757+5907 7
at M 1.8 1012M . In addition, our derived N is hours of HDS integration.
bulge ⊙ H
×
similar to the one Tremonti, Moustakas, Diamond-Stanic
9.2. Imaging of outflowing gas
(2007) derived for outflow in the massive post-starburst
galaxies (NH =2 1020 cm s−1). They also pointed out The blue component of [OIII] emission line has the
thatAGNsprobab×lyexistinthosepost-starburstgalaxies same outflow velocity to the trough at 1000 km s−1.
∼
because they get better fit to the spectra with a feature- Furthermore, the derived density and ionization param-
less blue power-law continuum and they detected high- eter values for the trough are typical for the narrow-line
excited emission-lines such as [OIII] and [NeV]. Thus, regions of AGNs (see, e.g., Groves, Dopita, Sutherland
the outflows seen in the massive post-starburst galaxies 2004). The scale of 3.7 kpc in AKARI J1757+5907 cor-
may be same phenomena as outflows in quasars. These responds to 0.′′55. The extended nebular gas associated
facts may be one of the indications that the outflow is a with the outflow can be observed by the [OIII] emission
commonphenomenonamongmassivegalaxies. Tremonti, line by HST or the [SIII] λ9533 emission line by integral
Moustakas, Diamond-Stanic (2007) had to make several filed spectroscopy coupled with adaptive optics observa-
assumptions to derive the physical quantities of the out- tions from large ground-based telescopes. Currently, the
flow because high resolution spectroscopy for such faint solidanglesubtendedbytheoutflowistheparameterwith
targets is difficult and non-detection of the important di- the largest uncertainty. Imaging the outflowing gas will
agnostic absorption lines from FeII (both ground and ex- directly determine this parameter.
cited states) and HeI*. In contrast, our high-resolution
9.3. HST Cosmic Origins Spectrograph (COS) observa-
spectroscopy of BAL outflows in quasars, can constrain
tions
the total hydrogen column density and the distance of
outflowing gas from the nucleus, which yield less model The near UV side of COS offers two important
dependent estimates for M˙ and E˙ . gains for the scientific investigation of the outflow in
out k
AKARI J1757+5907. First, using the G230L grating we
9. Scientific gains with additional observations can cover the full range of 1333–2800˚A (observed frame)
with resolution of 3000 and obtain data with S/N 30
9.1. Additional Subaru HDS data ∼ ∼
per resolution element using a total of six HST orbits.
Weexpectdatawithamuchhighersignal-to-noiseratio Such data will allow us to connect the low ionization ab-
usingSubaru/HDSundergoodweatherconditions. Ahalf sorption studied in this paper with the higher ionization
hour exposure should have a signal-to-noise ratio of 35 phase seen in SiIV, CIV, and NV. In addition, this spec-
per pixel at 3800 ˚A where FeII* λ2396 is shifted to. Our tralrangecovers3pairsofSiII/SiII*linesthatcandeter-
currentupperlimittoFeII*λ2612isderivedfromthedata mine number density considerably lower than is possible
with signal-to-noise ratio of 30 per pixel. FeII* λ2396 is with the lines from the FeII E=385 cm−1 lines (essen-
2.3 times stronger absorption than FeII* λ2612. With tially the critical density of the SiII* level is an order of
two hours (4 1800 sec) of exposure time, we expect five magnitude lower than that of the FeII E=385). In the
×
times more strict upper limit of the column density for unlikely event that we will detect SiII, but not SiII* ab-
the excitedstate ofFeII (385cm−1). Suchanupper limit sorption, which means that the gas density is substan-
would translate to a hydrogen number density of nH < tially below the criticaldensity for the SiII* level,we will
102.9 cm−3. For this value of nH and our best model at haveCII/CII*transitionscoveredthatcandeterminethe
thatdensity(logUH =-1.66andlogNH =21.32),wefind number densitydownto ne 10cm−3. Thesediagnostics
R > 5.9 kpc, E˙ > 1.2 1044 Ω ergs s−1 and M˙ > practically guarantee that w∼e will be able to determine
k 0.2 out
340 Ω M yr−1. It is×of course possible that we’ll be n and therefore the distance of the outflow, as long as
0.2 ⊙ e
abletodetectFeII*λ2396absorption,whichwillgiveusa R<100 kpc.
determination of ne and therefore measurements (instead InadditionwecantargetthestrongestpairofSiII/SiII*
of lower limits) for M˙ and E˙ . lines(1260˚A,1265˚A)forahigherresolution(R 20,000)
out k ∼
In this paper our analysis focused on the troughs for in order to obtain a fully resolved trough where we can
which we can measure FeII column density. The other obtain ne as a function of velocity for all 9 outflow com-
six troughs have only upper limits for FeII column den- ponents. With6HST orbitsusingthe COS225Mgrating
sity given the current signal-to-noise ratio. In order to we canobtainhighenoughS/Nto determinene for25-50
study the relationship between the troughs, and obtain resolved velocity points across the full width of the out-
the integrated values of M˙out and E˙k for the full outflow, flow for a large dynamical range in ne (100<ne <3000
it is important to measure FeII column density in all the cm−3). This will allow for a sensitive tomography of the
troughs. Combined with our other measurements, this outflow and precise determination of the distance to each
will constrain the ionization parameter and total hydro- kinematic component.
gen column density in each individual trough. We detect
HeI* absorption from 3 different lines in almost all the We are grateful to the staff of Subaru Telescope es-
troughs (figure 3). The current data at HeI* 2945,3188 pecially T.-S. Pyo, and A. Tajitsu, for their assistance
haveasignal-to-noise-ratioof65perpixel. Inordertoget during our HDS observations. We thank T. Hattori for
similar or better signal-to-noise ratio at FeII resonance taking FOCAS spectra kindly and providing to us. We
lines, 2383,2586, and 2600, we will need two additional also thank Jong-Hak Woo, T. Kawaguchi and the ref-
8 Aoki et al. [Vol. ,
eree for their helpful comments. This work was done Hamann,F.W.,Barlow, T.A.,Chaffee,F.C.,Foltz,C.B.,&
when K. A. and K. T. K. were staying at the Physics Weymann,R. J. 2001, ApJ, 550, 142
Department of Virginia Tech in summer 2010. K. A. H¨aring, N.,& Rix,H.-W. 2004, ApJL, 604, L89
and K. T. K. thank the staff there for hospitality and Hewett, P. C. & Foltz, C. B. 2003, AJ, 125, 1784
Hopkins, P. F., Hernquist, L., Cox, T. J., & Kerˇes, D. 2008,
support during their stay. We acknowledge support
ApJS, 175, 356
form NSF grant AST 0837880. This publication makes
Ishihara, D., et al. 2010, A&A,514, A1
use of data products from the Two Micron All Sky
Iye,M., et al. 2004, PASJ,56, 381
Survey, which is a joint project of the University of
Kashikawa, N., et al. 2002, PASJ, 54, 819
Massachusetts and the Infrared Processing and Analysis
Kawaguchi, T. 2003, ApJ,593, 69
Center/California Institute of Technology, funded by the Kellermann,K.I.,Sramek,R.,Schmidt,M., Shaffer,D.B., &
National Aeronautics and Space Administration and the Green, R. 1989, AJ, 98, 1195
National Science Foundation. Knigge, C., Scaring, S., Goad, M. R., & Cottis, C. E. 2008,
MNRAS,386, 1426
References Korista, K. T., Bautista, M., A., Arav, N., Moe, M.,
Costantini, E., & Benn, C. 2008, ApJ, 688, 108
Kraemer, S. B., Ruiz, J. R., & Crenshaw, D. M. 1998, ApJ,
Aoki, K., Kawaguchi, T., & Ohta, K. 2005, ApJ, 618, 601
508, 232
Arav, N., Brotherton, M. S., Becker, R. H., Gregg, M. D.,
Mathews, W. G. & Ferland, G. J. 1987, ApJ,323, 456
White, R.L., Price, T., & Hack,W. 2001, ApJ, 546, 140
Moe, M., Arav, N., Bautista, M., A., & Korista, K. T. 2009,
Arav, N., Kaastra, J., Kriss, G. A., Korista, K. T., Gabel, J.,
ApJ, 706, 525
& Proga, D.2005 ApJ,620, 665
Moll, R.,et al., 2007, A&A,463, 513
Arav, N., Moe, M., Costantini, E., Korista, K. T., Benn, C.,
Monet, D.,G., et al. 2003, AJ, 125, 984
& Ellison, S. 2008, ApJ, 681, 954
Murakami, H.,et al., 2007, PASJ,59, 369
Arav,N.,Dunn,J.P.,Korista,K.T.,Edmonds,D.,Gonz´alez-
Noguchi, K., et al. 2002, PASJ,54, 855
Serrano, J. I., Benn, C., & Jim´enez-Luj´an, F. 2010,
Olive, K. A. & Scully, S. T. 1996, International Journal of
submitting to ApJ
Modern Physics A, 11, 409
Ballero, S. K., Matteucci, F., Ciotti, L., Calura, F., &
Onken C. A., Ferrarese, L., Merritt, D., Peterson, B. M.,
Padovani, P. 2008, A&A,478, 335
Pogge, R. W., Vestergaard, M., & Wandel A. 2004, ApJ,
Barlow, T. A. & Sargent, W. L. W. 1997, AJ, 113, 136
615, 645
Bautista, M. A., Dunn, J. P., Arav, N., Korista, K. T.,
Schlegel,D.J.,Finkbeiner, D.P.,&Davis,M.1998, ApJ,500,
Moe, M., & Benn, C. 2010, ApJ,713, 25
525
Brotherton, M. S., Arav, N., Becker, R. H., Tran, H. D.,
Skrutskie,M. F., et al. 2006, AJ, 131, 1163
Gregg, M. D., White, R. L., Laurent-Muehleisen, S. A.,
Tremonti, C. A., Moustakas, J., & Diamond-Stanic, A. M.
& Hack, W. et al. 2001, ApJ, 546, 134
2007, ApJL,663, L77
Bennert, V. N., Treu, T., Woo, J.-H., Malkan, M. A.,
Trump, J. et al. 2006, ApJS,165,1
Auger, M. W., Gallagher, S., & Blandford, R. D. 2010,
VandenBerk, D.E. et al., 2001, AJ, 122, 549
ApJ, 708, 1507
Wampler,E.J.,Chugai,N.N.,&Petitjean,P.1995,ApJ,443,
Bentz,M.C.,Peterson,B.M.,Pogge,R.W.,Vestergaard,M.,
586
& Onken,C. A. 2006, ApJ, 644, 133
Weymann, R. J., Morris, S. L., Foltz, C. B., & Hewett, P. C.
Bru¨ggen, M. & Scannapieco, E. 2009, MNRAS,398, 548
1991, ApJ,373, 23
Ciotti, L., Ostriker, J. P., & Proga, D. 2010, ApJ,717, 708
Zhang, S., Wang, T.-G., Wang, H., Zhou, H., Dong, X.-B. &
Collins, N. R., Kraemer, S. B., Crenshaw, D. M., Ruiz, J.,
Wang, J.-G., ApJ, 714, 367
Deo, R., & Bruhweiler, F. C. 2005, ApJ, 619, 116
de Kool, M., Arav, N., Becker, R. H., Gregg, M. D.,
White, R. L., Laurent-Muehleisen, S. A., Price, T., &
Korista, K. T. 2001, ApJ,548, 609
Condon, J. J., Cotton, W. D., Greisen, E. W., Yin, Q. F.,
Perley, R. A., Taylor, G. B., & Broderick, J. J. 1998, AJ,
115, 1693
Dai, X., Shankar,F. and Sivakoff,G. R., arXiv:1004.0700
de Kool, M., Becker, R. H., Arav, N., Gregg, M. D.,
White, R.L. 2002, ApJ,570, 514
DiMatteo,T.,Springel,V.,&Hernquist,L.2005,Nature,433,
604
Dunn,J. P. et al. 2010, ApJ,709, 611
Ferland,G.J.,Korista,K.T.,Verner,D.A.,Ferguson,J.W.,
Kingdon, J. B., Verner,E. M. 1998, PASP,110, 761
Fuhr, J., & Wiese, W. 2006, J. Phys. Chem. Ref. Data, 35,
1669
Ganguly, R., & Brotherton, M. S.2008, ApJ,672, 102
Groves, B. A., Dopita, M. A., Sutherland, R. S. 2004, ApJS,
153, 9
Hall, P.B., et al. 2002, ApJS,141, 267
No. ] AKARI J1757+5907 9
Fig. 3. Outflow troughs in AKARI J1757+5907. Ordinate
isanormalizedfluxdensity,andabscissaisvelocityfromthe
Fig. 1. Low resolution spectrum of AKARI J1757+5907.
systemicredshift(z=0.61525). Blendedpartsaredenotedas
Ordinate is a flux density corrected for the Milky Way ex-
dottedspectra. Thedashedverticallineindicatestheposition
tinction in units of erg s−1 cm−2 ˚A−1, and abscissa is the
ofthebluecomponentof[OIII]emissionline. Thevelocityof
observed wavelength in angstrom. The rest wavelength is that[OIII]correspondstothetroughat∼−1000kms−1.
given along the top axis. The hatched regions indicate the
placeoftelluricabsorption.
Fig. 2. The Hβ-[OIII] region of AKARI J1757+5907. (a)
Continuum subtracted spectrum. The dotted line is the
FeII template produced from I Zw 1. The FeII tem-
plate is broadened and scaled at 5170 ˚A. Note that the
FeII template is clearly different from FeII emission line of
AKARIJ1757+5907at4600-4700˚Aand5260-5400˚A.(b)Fit
of the spectrum. Hβ, [OIII] doublet, and FeII λλ4925,5020
are fitted with a three Gaussians (cyan, blue and red lines),
twosets oftwo Gaussians (green lines),andtwo Lorentzians
(magentalines),respectively.
10 Aoki et al. [Vol. ,
Fig. 4. Top: AbsorptiontroughsfromHeIλ3890,3190,and Fig. 5. Similar to Figure 4. Top: This shows the trough
2946asred,green,andbluehistograms,respectively. Vertical profiles for the MgII doublet lines λλ2796, 2804 (blue and
barsrepresentthe statistical uncertainties inthe residualin- red histograms, respectively) and their associated statistical
tensity. Bottom: Velocity-dependent column density deter- uncertainties. Velocity-dependentcolumndensitydetermina-
minations for HeI* from the AOD (red histograms), C(v) tionsfromtheAOD(redhistograms),C(v)(bluehistograms),
(blue histograms), and power-law (green histograms) solu- andpower-law(greenhistograms)solutions. Thevelocity-in-
tions. The velocity-integrated HeI* column density values tegrated column density values for the three methods and
for the three methods and their associated statistical uncer- theirassociatedstatistical uncertainties arealsolisted.
taintiesarealsolisted.