A low-luminosity type-1 QSO sample: II. Tracing circumnuclear star formation in HE 1029-1831 with SINFONI
Gerold Busch, Semir Smajić, Julia Scharwächter, Andreas Eckart, Mónica Valencia-S., Lydia Moser, Bernd Husemann, Melanie Krips, Jens Zuther
AAstronomy & Astrophysics manuscript no. he1029 c (cid:13)
ESO 2018November 8, 2018
A low-luminosity type-1 QSO sample
II. Tracing circumnuclear star formation in HE 1029-1831 with SINFONI (cid:63)
Gerold Busch , Semir Smaji´c , , Julia Scharwächter , Andreas Eckart , , Mónica Valencia-S. , Lydia Moser , BerndHusemann , Melanie Krips , and Jens Zuther I. Physikalisches Institut, Universität zu Köln, Zülpicher Str. 77, 50937 Köln, Germanye-mail: [email protected] Max-Planck-Institut für Radioastronomie, Auf dem Hügel 69, 53121 Bonn, Germany LERMA (CNRS: UMR 8112), Observatoire de Paris, 61 Av. de l’Observatoire, 75014 Paris, France European Southern Observatory, Karl-Schwarzschild-Str. 2, 85748 Garching b. München, Germany Institut de Radio Astronomie Millimétrique (IRAM), 300 rue de la Piscine, Domaine Universitaire de Grenoble, 38406 St. Martind’Hères, FranceReceived;
ABSTRACT
Circumnuclear star formation and AGN feedback is believed to play a critical role in the context of galaxy evolution. The low-luminosity QSO (LLQSO) sample that contains 99 of the closest AGN with redshift z ≤ .
06 fills the gap between the local AGNpopulation and high-redshift QSOs that is essential to understand the AGN evolution with redshift. In this paper, we present the resultsof near-infrared H - and K - integral field spectroscopy of the inner kiloparsecs of the LLQSO HE 1029-1831 with SINFONI. Line mapsshow that ionized hydrogen gas is located in spiral arms within the stellar bar and a circumnuclear ring. Line fluxes and diagnosticline ratios indicate recent or ongoing star formation in the circumnuclear region and the presence of young and intermediate-agestellar populations in the bulge. In particular, we find traces of an intense starburst in the circumnuclear region that has begun around100 Myr ago but has declined to a fraction of the maximum intensity now. We estimate the dynamical bulge mass and find that thegalaxy follows published M BH − M bulge relations. However, bulge-disk decomposition of the K -band image with B udda reveals that HE1029-1831 does not follow the M BH − L bulge relations of inactive galaxies. We conclude that the deviation from M BH − L bulge relationsof inactive galaxies in this source is rather caused by young stellar populations and not by an undermassive black hole. Key words.
Galaxies: active – Galaxies: Seyfert – Galaxies: individual: HE 1029–1831
1. Introduction
Tight correlations between the mass of the central supermassiveblack hole (SMBH) and properties of the hosting galaxy, partic-ularly its central bulge component, have been found (Magorrianet al. 1998; Ferrarese & Merritt 2000; Marconi & Hunt 2003;Häring & Rix 2004; Graham & Driver 2007; Savorgnan et al.2013) and are seen as evidence for a co-evolution of black holesand their host galaxies. More recent studies have revealed galax-ies that do not follow these relations: Kormendy et al. (2011)discuss that only “classical bulges” that formed by galaxy merg-ers show correlations with the SMBH while disk-like “pseudob-ulges” that are forming by secular evolution do not. Other au-thors suggest that di ff erent formation processes (gaseous forma-tion processes vs. dry merging) might become visible in di ff er-ing SMBH-bulge scaling relations (Graham & Scott 2013; Scottet al. 2013). Further, Läsker et al. (2014) claim that, providinga thorough decomposition, the correlation between black holemass and total host galaxy luminosity could be equally tight asthe correlation between black hole mass and bulge luminosity.In our recent study (Busch et al. 2014), we probed the M BH − L bulge relation of low-luminosity type-1 QSOs (LLQSOs).We performed a careful decomposition of K -band images withthe B udda -code (de Souza et al. 2004; Gadotti 2008) and find (cid:63) Based on observations with ESO-VLT, proposal no. 093.B-0718(A) that the bulges do not follow the published M BH − L bulge re-lations for inactive galaxies of Marconi & Hunt (2003); Vikaet al. (2012); Graham & Scott (2013); Kormendy & Ho (2013).A deviation of type-1 AGN from the SMBH-bulge relations isin agreement with previous studies in the optical (Nelson et al.2004; Kim et al. 2008; Bennert et al. 2011). Assuming that aSMBH-bulge relation traces a co-evolution, this could hint at adi ff erent evolution of inactive and active galaxies. On the otherhand, a deviation from the relations of inactive galaxies couldhint at di ff erent properties of active galaxies compared to inac-tive ones, caused by direct interaction through AGN feedback.Particularly, we discussed an undermassive black hole that isgrowing “towards” the relation or an overluminosity of the bulgeproduced by young and / or intermediate-age stellar populations.With the emergence of AO-assisted integral-field spectroscopy,many studies have been focused on the circumnuclear star for-mation properties in AGN (e.g., Zuther et al. 2007; Davies et al.2007; Böker et al. 2008; Ri ff el et al. 2010). However, the impacton the bulge and overall host galaxy and the interaction betweencentral star formation and AGN feedback is still under discus-sion. Article number, page 1 of 16 a r X i v : . [ a s t r o - ph . GA ] D ec & A proofs: manuscript no. he1029
The low-luminosity type-1 QSO sample (LLQSO sample) is asubsample of the Hamburg / ESO survey for bright UV-excessQSOs (Wisotzki et al. 2000) that contains only the 99 bright-est and nearest QSOs with a redshift z ≤ .
06. In many prop-erties, e.g. redshift, gas masses, and luminosities, the LLQSO-sample lies between the NUGA sample (NUclei of GAlaxies;e.g., García-Burillo et al. 2003; Combes et al. 2004; García-Burillo et al. 2005; Krips et al. 2007a) and the Palomar Green(PG) QSO sample. This connects the advantages of cosmologi-cal proximity and of higher spatial coverage. The first enables usto observe the processes in the centers of the galaxies in detailwith high spatial resolution. The latter raises the probability ofobserving QSOs with high luminosity and accretion rate (Kripset al. 2007b; Bertram et al. 2007; Moser et al. 2012).In Table 1, we compare the source HE 1029-1831 with NGC3227, the nearest Seyfert-1 galaxy, and 3C 273, the nearestquasar. We see that the LLQSO has a much higher bolometric lu-minosity and accretion rate (traced by the Eddington ratio) thanthe nearby Seyfert galaxy. However, it is much closer than thenext quasar, the resolution being better by a factor of four. Thus,LLQSOs form an ideal sample to study active galaxies with sig-nificant accretion activity at a resolution that is still good enoughfor detailed analysis.Bertram et al. (2007) and Krips et al. (2007b) observed 39galaxies from the LLQSO sample in CO from which 27 havebeen detected. They report that most galaxies are rich in molec-ular gas, with a range of gas masses from 0 . × M (cid:12) to9 . × M (cid:12) . Furthermore, they find that the gas mass is cor-related with the star formation and AGN activity as indicated bythe far-infrared luminosity. König et al. (2009) searched for 21cm H i emission in these 27 CO-detected galaxies with the Ef-felsberg 100m-telescope. They find neutral atomic gas massesranging from 1 . × M (cid:12) to 3 . × M (cid:12) but no strong cor-relation between gas mass and infrared emission that traces starformation and AGN activity. Fischer et al. (2006) observed ninegalaxies from the sample with the near-infrared spectrographISAAC (VLT) and find that the spectra show signs for both,AGN activity but also star formation. In our near-infrared study(Busch et al. 2012, 2014) we find that the LLQSOs do not followpublished M BH − L bulge relations for inactive galaxies. A reasoncould be massive star formation in the circumnuclear regions.Integral-field spectroscopy with SINFONI is the preferred toolto examine the excitation mechanisms of emission lines as wellas the kinematics of gas and stars in a spatially resolved wayand learn about the properties of star formation and AGN in thecenters of LLQSOs.Since HE 1029-1831 shows strong CO(1-0) emission (Kripset al. 2007b, corresponding to a cold gas mass of ∼ . × M (cid:12) ), it is an ideal object for a pilot study to demonstratethat with SINFONI we can successfully trace star formation re-gions as well as hot molecular and ionized gas reservoirs alsoin LLQSOs which are more distant than most of the previouslyanalyzed (e.g. by the NUGA team) objects. Fischer et al. (2006) analyze NIR
JHK -images as well as H + K -longslit spectra of nine LLQSOs, including HE 1029-1831,obtained with ISAAC at the VLT. In Fig. 1, we see that HE 1029-1831 is a barred spiral galaxy with two prominent spiral arms.The bar has a length of about 8 (cid:48)(cid:48) ≈ . Fig. 1.
ISAAC K -band image of HE1029-1831. The 8 (cid:48)(cid:48) × (cid:48)(cid:48) and 3 (cid:48)(cid:48) × (cid:48)(cid:48) FOVs of our SINFONI observations are marked. The image has beensmoothed with a 3-pixel Gaussian and cleared from bad pixels.
In their spectra, they can separate the Pa α line into broadand narrow component; the broad component has a full width athalf maximum (FWHM) of 2081 km s − . Furthermore, they findextended molecular hydrogen emission. Due to the FWHM(Pa α )of about 2000 km s − , they classify HE 1029-1831 as narrow-line Seyfert-1 galaxy (NLSy1). This is supported by Rodríguez-Ardila et al. (2000), using the name “CTS J04.08”, who find aFWHM(H α ) of 1870 km s − (but see discussion about NLSy1sin Valencia-S. et al. 2012b).Krips et al. (2007b) observed HE 1029-1831 in the CO(1-0) and CO(2-1) line emission with the IRAM Plateau de BureInterferometer and obtain a total gas mass of M (H + He) ≈ × M (cid:12) . For the spatial extent of the emission, they find a FWHMof ≈ (7 ± (cid:48)(cid:48) , which means that the observed CO emission comesfrom a region including the complete stellar bar. The H i masshas been measured from observations with the E ff elsberg 100m-telescope and is (6 . ± . × M (cid:12) (König et al. 2009).From optical line ratios, HE 1029-1831 is classified asSeyfert 1.5 galaxy in the 12th catalog of quasars and AGN byVéron-Cetty & Véron (2006). Kewley et al. (2001) took opticalspectra and determine the position of the galaxy in optical diag-nostic diagrams (using the name IRAS 10295-1831). Dependingon the diagram, the galaxy is classified as AGN, starburst Article number, page 2 of 16erold Busch et al.: LLQSO sample - Tracing circumnuclear star formation in HE 1029-1831 with SINFONI
Table 1.
Comparison of the nearest Seyfert-1 galaxy NGC 3227, our source HE 1029-1831, and the nearest quasar 3C 273. source redshift scale resolution (*) log( L bol [W]) log( M BH [ M (cid:12) ]) λ Edd
NGC 3227 0.004 100 pc / (cid:48)(cid:48)
13 pc 36.86 7.64 0.013HE 1029-1831 0.0404 800 pc / (cid:48)(cid:48)
110 pc 37.6 7.2 0.33C 273 0.158 3150 pc / (cid:48)(cid:48)
410 pc 39.94 8.95 0.8
Notes. ( * ) Assuming a resolution of 0 (cid:48)(cid:48) .
13 as we reach with AO correction in our observations with the 3 (cid:48)(cid:48) × (cid:48)(cid:48) FOV. The redshift is taken from theNED database. Black hole mass and bolometric luminosity of NGC 3227 and 3C 273 are from Woo & Urry (2002) and references therein. or composite. From this, we expect that in HE 1029-1831contributions from both, AGN and starburst, are important.In this paper, we present near-infrared integral-field spec-troscopy of HE 1029-1831 with SINFONI. HE 1029-1831 isthe first of a set of galaxies from the LLQSO sample that hasbeen observed with SINFONI. From the 3D-spectroscopy, weaim to learn about the stellar and gas masses as well as kinemat-ics in LLQSOs. We will determine reliable mass-to-light ratiosand trace the star formation history in the centers of type-1 AGN.The paper is organized as follows: In Sect. 2, we presentthe observations with SINFONI as well as preceding observa-tions with ISAAC (near-infrared image) and describe our reduc-tion procedures. In Sect. 3, we describe our methods, that is thebulge-disk decomposition with B udda , emission line fitting, andstellar continuum subtraction with pPXF, and present the results.In Sect. 4, we discuss the results, particularly we estimate blackhole masses, determine gas masses and gas excitation mecha-nisms, and discuss the impact of star formation. In Sect. 5, wegive a short summary and conclusions.Throughout this paper, we use a standard cosmology with H =
70 km s − Mpc − , Ω m = .
3, and Ω Λ = .
7. At a redshiftof z = . D L =
177 Mpc and a scale of 0 . / (cid:48)(cid:48) .
2. Observation and data reduction
We present results of high-resolution NIR observations of HE1029-1831 that have been carried out April 20-22, 2014 withSINFONI at the Unit Telescope 4 of the ESO Very Large Tele-scope in Chile (VLT, Eisenhauer et al. 2003; Bonnet et al. 2004).We started the observations with the 0 (cid:48)(cid:48) .
25 plate scale, cor-responding to a field-of-view (FOV) of 8 (cid:48)(cid:48) × (cid:48)(cid:48) without adap-tive optics. Then, we used the 0 (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) . We used the H + K grating provid-ing a spectral resolution of R H + K = ff setsof ± (cid:48)(cid:48) for the 8 (cid:48)(cid:48) × (cid:48)(cid:48) -FOV and ± (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) -FOV wasused to minimize the impact of dead pixels. The total integrationtime on source was 4500 s for the large and 3000 s for the smallFOV, with additional 2250 s and 1500 s for sky observations. Inbetween, telluric standards have been observed.In the raw files, detector specific problems occured that wecorrected following the procedure described in the appendix ofSmaji´c et al. (2014). Then, we used the SINFONI pipeline fordata reduction up to single-exposure-cube reconstruction. Thealignment and final coaddition of the single-exposure-cubes hasbeen performed with our own dpuser routines. http: // / ˜ott / dpuser / dpuser.html For the telluric correction, we used G2V stars that we haveobserved in between. A high-quality solar spectrum (Maiolinoet al. 1996) that has been convolved down to the resolution ofthe telluric spectrum was used to correct the features of the G2Vstar. At the spectral edges and in the spectral region between H -and K -band, the solar spectrum had to be interpolated by a blackbody function with T = (cid:48)(cid:48) . − (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) , which corresponds to linear scales of 320 − (cid:48)(cid:48) × (cid:48)(cid:48) , we measure a FWHM of 0 (cid:48)(cid:48) . λ . µ m were taken as areference for the science target. The magnitude of the standardstar was taken from the 2MASS point source catalog (Skrutskieet al. 2006). In order to convert the 2MASS magnitude into aflux density, we use the Spitzer Science Center Magnitude toFlux Density converter . K -band imaging A K -band image of HE1029-1831, observed with ISAAC(Moorwood et al. 1998) at the VLT on April 20, 2003, is avail-able in the ESO-archive. ISAAC provides a FOV of 152 (cid:48)(cid:48) × (cid:48)(cid:48) with a pixelscale of 0 . (cid:48)(cid:48) / pixel. Four frames have been ob-served in a TSST jitter pattern with an exposure time of 2 seceach. After subtracting the sky frames, the two frames have beenshifted to a common reference frame and a mean image was cre-ated. Calibration was done using 2MASS magnitudes of fore-ground stars. For details on the data, see Fischer et al. (2006).
3. Results
Parametric modeling of galaxy images is an important tool todisentangle the light of di ff erent galaxy components, such asbulge, disk, bar, AGN, to measure their structural parameters anddetermine light fractions of the galactic components. These val-ues are essential for estimates of stellar masses and are used toestablish relations e.g. between black hole mass and bulge lumi-nosity.In this work, we use B udda (Bulge / Disk DecompositionAnalysis; de Souza et al. 2004; Gadotti 2008) for the two-dimensional decomposition of the K -band image. http: // ssc.spitzer.caltech.edu / warmmission / propkit / pet / magtojy / http: // / ˜dgadotti / budda.htmlArticle number, page 3 of 16 & A proofs: manuscript no. he1029
To model the disk component, we use an exponential func-tion µ disk ( r ) = µ + . rh r (1)with central surface density µ and the characteristic scale-length h r (Freeman 1970).Bulge and bar component are modeled by a Sérsic function µ ( r ) = µ e + c n (cid:32) rr e (cid:33) / n − (2)with an e ff ective radius r e that is the radius that includes halfof the light emitted by the respective component (“half-light ra-dius”) and the e ff ective surface brightness µ e ≡ µ ( r e ). n is theSérsic index and c n = . . n − . n (Sersic 1968; Caon et al. 1993).The AGN component is modeled by a Mo ff at function withwidth fixed to that of the seeing PSF (measured by fitting Gaus-sian functions to point sources). The width is therefore kept fixedand only the peak intensity value is varied. B udda calculatesthe minimum valid count as the square root of the noise level.Here, the limiting surface brightness is µ K = . − .All pixels with count level below will be masked. Busch et al.(2014) give more details on the fitting process and limits of 2D-decomposition.Figure 2 shows the results of the 2D-decomposition of HE1029-1831: In the upper row, from left to right the original K -band image, the B udda model image and a residual image. In thelower row, we show radial profiles that have been extracted usingthe E llipse task of I raf . In the fits of the model components,ellipticity and position angle have been fixed to the values of theoriginal image.The fit results in a bulge component with an e ff ective ra-dius of r e = . (cid:48)(cid:48) ( = .
47 kpc) and an apparent magnitude of m bulge = .
69 (absolute magnitude: M bulge = − . h r = . (cid:48)(cid:48) ( = .
16 kpc) and a magnitudeof m disk = .
04 ( M disk = − . n = .
2. A Sérsic indexbelow 2 is indicative for a disk-like pseudobulge (Kormendy& Kennicutt 2004). However, on the one hand the Sérsic indexshould not be used as a sole indicator (Gadotti 2009) and on theother hand, since the central light distribution is heavily a ff ectedby the nuclear point source, it is almost impossible to determinethe Sérsic index reliably. The spectra obtained with SINFONI show a big variety of emis-sion lines. In the top right panel of Fig. 3, we mark three spots,the nuclear region and two o ff -nuclear regions with a radius of0 (cid:48)(cid:48) .
15 each. Furthermore, we extract one spectrum from an aper-ture 2 (cid:48)(cid:48) south of the nucleus. Here we chose a larger aperturewith a radius of 0 . (cid:48)(cid:48) matching the lower spatial resolution in the8 (cid:48)(cid:48) × (cid:48)(cid:48) data set. The maps show the continuum emission around2 . µ m from the SINFONI-cube, left for the 8 (cid:48)(cid:48) × (cid:48)(cid:48) and right forthe 3 (cid:48)(cid:48) × (cid:48)(cid:48) FOV. In the lower panels, we show three spectra thatare extracted from these regions. In the nuclear spectrum, broadline components and the coronal line [Si vi ] are apparent, whichindicate the presence of strong radiation from the central AGN.In the o ff -nuclear spectra, stellar features like the CO band headsbecome more prominent. Moreover, 3D-spectroscopy enables us to extract spectra and fit the emission lines in every spectral pixel(spaxel), which results in maps of the emission line fluxes.All emission lines have been fitted with a single Gaussiancomponent plus zero-order polynomial. In case of the hydro-gen recombination lines Pa α and Br γ , a second component, ac-counting for the light emitted from the broad line region (BLR)was added. Emission lines that are close to each other (e.g.,[Si vi ] and H λ µ m) have been fitted simultaneously. Exam-ples for these fits are displayed in Fig. 4. We use the python -implementation of mpfitexpr (Markwardt 2009). Errors of the fitparameters are calculated from the covariance matrix. For testingpurposes, we also determined uncertainties by generating 100Monte Carlo realizations of the input spectrum, adding Gaussiannoise with a width corresponding to the standard deviation in aline-free region. Both methods deliver comparable results. If notnoted otherwise, we clipped values in the maps with a relativeerror larger than 30%.In Table 2, we present the flux measurements for the emis-sion lines in three apertures of radius 0 (cid:48)(cid:48) .
15 each that are markedin the continuum image in Fig. 3.Two strong hydrogen recombination lines were detected: Br γ and Pa α . In the central region, both show two spectral compo-nents, a broad and a narrow one. The broad components show aFWHM of ≈ − , indicative for a type-1 AGN. Accord-ing to the unified model (e.g., Antonucci 1993; Urry & Padovani1995), this emission stems from clouds very close ( r ∼ lightdays,see review of Peterson 1993) to the central supermassive blackhole which are moving at high velocities.HE 1029-1831 shows extended narrow emission in Br γ andPa α . In the Pa α emission two spiral arms in the ionized gas emis-sion are clearly visible (Fig. 5). Since Br γ is less strong, thearms are not as prominent (Fig. A.1). The AO-assisted 3 (cid:48)(cid:48) × (cid:48)(cid:48) -observations reveal a patchy ring in ionized gas emission ofabout 0 (cid:48)(cid:48) . =
240 pc in radius and an ellipticity of (cid:15) = . ii ]-emission does not peak in the center but in twoo ff -nuclear regions that are identical with the two o ff -nuclearPa α / Br γ peaks (Fig. 6). Thus, we conclude that the strong cir-cumnuclear gas emission does not mainly stem from the narrow-line region but traces a circumnuclear star-formation region.Furthermore, we see emission from the coronal line [Si vi ].We fit a Gaussian function to the spatial distribution and get aFWHM of 0 (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) -FOV and 0 (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) FOV. Coronal lines are forbidden transitions in the ground lev-els of highly ionized atoms that have ionization potentials above100 eV. Thus, they are free of contribution from star formationand directly associated with the AGN. Müller Sánchez et al.(2006) argue that coronal line emission stems from the innerNLR and that they are probably associated with outflows. Re-cently, Müller-Sánchez et al. (2011) and Mazzalay et al. (2013a)observed extended emission in the coronal lines of NGC 1068as well as a great complexity in the morphology and kinematics.This indicates that the nature of coronal line emission is far frombeing understood.For the broad components of Pa α and Br γ , we find a FWHMof 0 (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) -FOV and 0 (cid:48)(cid:48) . − (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) -FOV.For the 8 (cid:48)(cid:48) × (cid:48)(cid:48) FOV, these values agree with the values derivedfrom the telluric star observations in Sect. 2.1. For the 3 (cid:48)(cid:48) × (cid:48)(cid:48) -FOV, the values derived here are slightly higher. We claim thatthey are a more robust method to derive the spatial resolutionsince they are extracted from the analyzed data itself. Article number, page 4 of 16erold Busch et al.: LLQSO sample - Tracing circumnuclear star formation in HE 1029-1831 with SINFONI µ K ( m a g / a r c s e c ) bulgediskbaragnmodel g a l - m o d E ll . P . A .
10 0 10x offset (arcsec)10010 y o ff s e t ( a r c s e c ) g a l a xy ( K - m a g )
10 0 10x offset (arcsec) m o d e l ( K - m a g )
10 0 10x offset (arcsec) r e s i d uu m ( K - m a g ) Fig. 2.
Decomposition of HE 1029-1831 with B udda . We show ( from left to right ) the original K -band ISAAC-image, the B udda -model and theresiduum obtained by division galaxy / model. Here, blue indicates that the original image has more flux than the model while red indicates that themodel has more flux than the original image. In the lower row , we show an elliptically averaged radial profile of the galaxy, the single componentsand the model. Next to it, we show the di ff erence galaxy-model, ellipticity and position angle. Several H molecular hydrogen emission lines are detected:H (1-0)S(3) λ µ m, H (1-0)S(2) λ µ m, H (1-0)S(1) λ µ m, H (1-0)S(0) λ µ m, and H (2-1)S(1) λ µ m.Other lines that are visible are He i as well as several CO stellarabsorption bands. Under the assumption that the central bulge component is sup-ported by random motion rather than ordered rotation, the stel-lar velocity dispersion can be used for virial mass estimatesand black hole estimates via the M BH − σ ∗ relation. To mea-sure the stellar velocity dispersion σ ∗ , we fit stellar templates tothe K -band spectrum that we integrated over an aperture corre-sponding to the e ff ective radius measured in Sect. 3.1. We usethe penalized pixel-fitting algorithm (pPXF, Cappellari & Em-sellem 2004) and take stellar templates from the Gemini Spec-tral Library of Near-IR Late-Type Stellar Templates (Winge et al.2009) which consists of 11 giant and supergiant stars with spec-tral classes from G8 to M5 that have been observed with theGEMINI integral-field spectrograph NIFS. The H + K grating ofSINFONI has a spectral resolution of R = λ = . µ m which we confirmed by measuring thewidth of OH-lines. The stellar templates have a di ff erent spectral resolution. Therefore, the templates had to be degraded down tothe resolution of SINFONI first. Regions that contain emissionlines or telluric features are masked. Fig. 7 shows the best fit witha stellar velocity dispersion of σ ∗ = (96 ±
12) km s − . To test ourresults, we also use a set of theoretical spectra of red giant andsupergiant stars from Lançon et al. (2007) and get a consistentresult σ ∗ = (111 ±
14) km s − . In the following sections, we willuse a mean value of σ ∗ =
104 km s − . To estimate the uncertain-ties, we perform a Monte Carlo simulation. We generate 1000spectra by adding Gaussian noise to the original spectrum. Thewidth of the normal distribution of the noise is given by the stan-dard deviation of the flux in a line-free region. The uncertaintyis then given by the standard deviation of the 1000 fit results. Extinction can have a significant e ff ect on the line fluxes. There-fore, we use hydrogen recombination line ratios to estimate theextinction. We follow the extinction law derived by Cardelli et al.(1989) to estimate the extinction in the near-infrared: A V = . × log (cid:32) ( f Br γ / f Pa α ) obs ( f Br γ / f Pa α ) intr (cid:33) (3) Article number, page 5 of 16 & A proofs: manuscript no. he1029 ∆ RA [arcsec]432101234 ∆ D E C [ a r c s e c ] spot 3 ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] nucleusspot 1spot 2 C O [ F e II ] P a α H [ S i V I ] H H e I H B r γ H H C O C O nucleus f λ [ − W m − µ m − ] spot 1 spot 2 wavelength λ [ µm ] spot 3 Fig. 3.
In the upper panels we show the continuum emission around 2 . µ m in the 8 (cid:48)(cid:48) × (cid:48)(cid:48) and the 3 (cid:48)(cid:48) × (cid:48)(cid:48) FOV, in unit 10 − W m − µ m − . Inthe bottom panels , spectra extracted from three apertures with radius 0 (cid:48)(cid:48) .
15 in the 3 (cid:48)(cid:48) × (cid:48)(cid:48) -FOV and from one aperture with radius 0 (cid:48)(cid:48) . (cid:48)(cid:48) × (cid:48)(cid:48) -FOV that are marked in the continuum images are shown. The emission lines and the CO absorption band heads are marked in the spectra. where ( f Br γ / f Pa α ) obs is the observed line ratio between Br γ andPa α that is divided by the intrinsic line ratio ( f Br γ / f Pa α ) intr = .
083 taken from Osterbrock & Ferland (2006) assuming case Brecombination with a typical electron density of n e = cm − and a temperature T = K. In the nuclear aperture, we mea-sure an extinction of A V = . A V = . A V = − . A V are up to 100%.Thus, we decide not to correct the near-infrared fluxes for ex-tinction. However, we have to keep in mind that this introduces an error. In particular, ignoring an extinction of A V =
4. Discussion
The observed correlations between black hole mass and hostgalaxy properties imply that the evolution of the central super-massive black holes and their hosts might be connected to eachother. Black hole masses are thus of great interest. From sin-
Article number, page 6 of 16erold Busch et al.: LLQSO sample - Tracing circumnuclear star formation in HE 1029-1831 with SINFONI
Table 2.
Emission-line fluxes for three positions, the center and two o ff -center spots. Apertures have a radius of 0 (cid:48)(cid:48) . nucleus spot 1 spot 2line λ [ µ m] Flux FWHM Flux FWHM Flux FWHM[Fe ii ] 1.64400 26 ± . ± . . ± . α (narrow) 1.87561 304 ±
13 290 442 ± ± α (broad) 1.87561 920 ±
39 2150 * ±
15 2150 92 ±
10 2150 * H (1-0)S(3) 1.95756 20 . ± . . ± . . ± . vi ] 1.9634 43 ± . ± . . ± . (1-0)S(2) 2.03376 5 . ± . . ± . . ± . i ± . ± . . ± . (1-0)S(1) 2.12183 12 . ± . . ± . . ± . γ (narrow) 2.16612 26 ± . ± . . ± . γ (broad) 2.16612 135 ±
11 2150 * ± ± * H (1-0)S(0) 2.22329 5 ± . ± . . ± . (2-1)S(1) 2.24772 — — 4 . ± . . ± . Notes.
All flux measurements are in units of 10 − W m − . FWHMs are in km s − and not corrected for instrumental broadening. ( * ) Broad linewidth was fixed to the line width measured for the broad Pa α component in the nuclear aperture. λ [ µm ]1.01.52.02.53.03.54.0 f λ [ − W m − µ m − ] Pa α complex λ [ µm ]1.651.701.751.801.851.901.95 f λ [ − W m − µ m − ] Br γ and H complex λ [ µm ]1.651.701.751.801.851.90 f λ [ − W m − µ m − ] H and [SiVI] complex Fig. 4.
Three examples for the fits applied to the cubes. The plots show the fits of the corresponding complex in the nuclear aperture. gle epoch spectroscopy, we expect a black hole mass of aroundlog( M BH / M (cid:12) ) = . − .
4. We compare this mass to other massestimates (Table 3) and find an Eddington ratio of λ ≈ .
3, im-plying that the black hole is e ffi ciently growing.For 74 galaxies from the LLQSO sample, BH masses areavailable from Schulze et al. (2009), Schulze & Wisotzki (2010),and Schulze (priv. comm.). From single epoch spectroscopy,they compute BH masses using the scaling relation betweenbroad line region (BLR) size and continuum luminosity (Bentzet al. 2009) M BH = . f (cid:18) L W (cid:19) . (cid:18) σ H β km s − (cid:19) M (cid:12) . (4)We use their measurements for L and σ H β and apply therecent scale factor f = . M BH / M (cid:12) ) = . (Husemannet al. 2013, 2014) in order to deblend QSO and host galaxy emis-sion in optical IFS-datacubes obtained with WIFES (Dopita et al.2007, 2010) and estimate a black hole mass of log( M BH / M (cid:12) ) = . L and the FHWM of the broad H β component.We derive the black hole mass from the broad Pa α compo-nent following Kim et al. (2010) M BH = . ± . (cid:18) L Pa α W (cid:19) . ± . (cid:32) FWHM Pa α km s − (cid:33) M (cid:12) . (5)We measure a luminosity in the broad Pa α component of L Pa α, broad = . × W. With a broad line width of FWHM = − , we get a black hole mass of log M BH = .
3. How-ever, because of the uncertainties discussed in Sect. 3.4, we didnot correct for extinction. Extinction correction would result ina higher estimate.The stellar velocity dispersion, derived from the stellar con-tinuum fit in Sect. 3.3, can be used to estimate the black holemass from the M BH − σ ∗ relation. From the velocity disper-sion σ ∗ =
104 km s − , we find the following estimates forlog( M BH / M (cid:12) ): 7.0, 6.7, and 7.4 (using the relations of Gültekinet al. 2009; Graham et al. 2011; Kormendy & Ho 2013). Re-cently, Ri ff el et al. (2014a) reported that the stellar velocity dis-persion in spiral galaxies derived from CO band heads in thenear-infrared ( σ CO ) and from the Calcium triplet in the optical( σ opt ) di ff er from each other significantly. In the above men-tioned M BH − σ relations, the optical velocity dispersion is used.For HE 1029–1831, we correct our near-infrared velocity dis-persion using their best fit and derive a corresponding velocitydispersion of σ opt =
134 km s − in the optical. This results inBH masses that are higher by a factor of ≈ . M BH − σ relation are only esti-mations. The uncertainty will be high, at least 0 . M bulge = − .
57 (Sect. 3.1), wecan estimate the black hole mass that would be expected from M BH − L bulge relations for inactive galaxies. The mass estimates Article number, page 7 of 16 & A proofs: manuscript no. he1029 ∆ RA [arcsec]432101234 ∆ D E C [ a r c s e c ] Pa α - Flux F l u x [ − W m − ] ∆ RA [arcsec]432101234 ∆ D E C [ a r c s e c ] Pa α - Equivalent width E W [] ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] F l u x [ − W m − ] ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] E W [] Fig. 5.
From left to right:
Flux and equivalent width of the narrow component of the hydrogen recombination line Pa α . Upper row:
FOV 8 (cid:48)(cid:48) × (cid:48)(cid:48) , lower row: FOV 3 (cid:48)(cid:48) × (cid:48)(cid:48) . The center of the continuum emission is marked by a × . ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] nucleusspot 1spot 2 H (1-0)S(3) - Flux F l u x [ − W m − ] ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] nucleusspot 1spot 2 H (1-0)S(1) - Flux F l u x [ − W m − ] ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] nucleusspot 1spot 2 [FeII] - Flux F l u x [ − W m − ] Fig. 6.
Flux of ( from left to right ) H (1-0)S(3) λ µ m, H (1-0)S(1) λ µ m, and [Fe ii ] λ µ m. The center of the continuum emission ismarked by a × . The three apertures nucleus, spot1, and spot2 are marked as well. log( M BH / M (cid:12) ) range from 8.0 to 8.4 (following Marconi & Hunt2003; Vika et al. 2012; Graham & Scott 2013; Kormendy & Ho2013). As observed previously for other LLQSOs, the measuredblack hole masses are lower by a factor of ≈ M BH − L bulge relations of inactive galaxies.As discussed in Busch et al. (2014), this could be caused by over-luminous host galaxies that contain young stellar populations or by black holes that are in a growing phase and still undermassivein comparison to their inactive counterparts.The Eddington ratio is a measure of the accretion e ffi ciencyand is defined as λ ≡ L bol / L Edd . The bolometric luminosity canbe derived from the X-ray luminosity. The soft X-ray flux fromROSAT has been measured to be f soft X − ray = (2 . ± . × − W m − (Mahony et al. 2010), corresponding to a luminos- Article number, page 8 of 16erold Busch et al.: LLQSO sample - Tracing circumnuclear star formation in HE 1029-1831 with SINFONI µ m]0.70.80.91.01.11.21.3 f λ [ a r b i t r a r y un i t s ] σ ∗ =(96 ± km/s Fig. 7.
Fit of the stellar kinematics of the K -band spectrum of HE 1029-1831. The radiusof the aperture is r = . (cid:48)(cid:48) =
470 pc, cor-responding to the e ff ective radius. The spec-trum is shown in black, the best fitting combina-tion of NIFS templates in red. Spectral regionsthat are contaminated by emission lines or tel-luric features are masked and are marked by theshaded areas. The residuum, shifted verticallyfor clarity, is shown in gray. Table 3.
Black hole mass estimates from di ff erent relations. Used relation log( M BH / M (cid:12) ) broad line width A. Schulze (priv. comm.) 7.4J. Scharwächter (priv. comm.) 7.3Kim et al. (2010) 7.3 M BH − σ ∗ relation Gültekin et al. (2009) 7.0 ( + a )Graham et al. (2011) 6.7 b ( + a )Kormendy & Ho (2013) 7.4 ( + a ) M BH − L bulge relation c Marconi & Hunt (2003) 8.0Vika et al. (2012) 8.2Graham & Scott (2013) 8.4Kormendy & Ho (2013) 8.4 M BH − M bulge relation Häring & Rix (2004) 6.0-6.8Sani et al. (2011) 6.6-7.1Scott et al. (2013) 5.1-6.6 / d Kormendy & Ho (2013) 6.4-7.2
Notes.
Note that the derivation of uncertainties for black hole massesis complex since it includes measurement uncertainties (e.g. of the fluxmeasurement) but also the intrinsic scatter of the relations. Usually, theuncertainty will be at least 0.5 dex. (a) correction for the di ff erence be-tween near-infrared ( σ CO ) and optical ( σ opt ) stellar velocity dispersionaccording to Ri ff el et al. (2014a). ( b ) Using the relation for barred galax-ies. ( c ) Note that the M BH − L bulge relation is not a good estimator foractive galaxies according to Busch et al. (2014). ( d ) Using the relationfor core-Sérsic galaxies or Sérsic galaxies respectively. ity of L soft X − ray = (1 . ± . × W. With a conversion factorof ≈
50 (Hopkins et al. 2007), we get a bolometric luminosity of L bol , X − ray = (6 ± × W.Another possibility is to derive the bolometric luminos-ity from the AGN continuum luminosity at 5100 Å. Schar-wächter et al. (in prep.) subtract the host galaxy contribu-tion using QDeblend and derive a flux of F = . × − erg s − cm − Å − . Using L bol = f L × L with f L = L bol = . × W.The Eddington luminosity is given by L Edd (cid:27) . × ( M BH / M (cid:12) ) W. Thus, the range of black hole mass estimatesof log( M BH ) = (6 . − .
4) corresponds to Eddington luminosi-ties L Edd = (6 − × W. With the bolometric luminosities L bol = (2 − × W, we can estimate the Eddington ratio tobe around λ = (0 . − . λ = .
3. Although the range is very broad, we can point out that even the lowestpossible Eddington ratio is high, indicating an actively accret-ing black hole, i.e. a central black hole in a phase of substantialgrowth.
In our spectra, we observe a variety of emission lines, mostof them originate from ionized or molecular hydrogen but also[Fe ii ]. In recent 3D spectroscopy studies, it has been shown thatthe H and the [Fe ii ] emitting gas often have di ff erent flux distri-butions and kinematics. We find that [Fe ii ] and hydrogen recom-bination lines (Pa α and Br γ ) show similar distribution while H is more centrally distributed (see Figs. 5 and 6). The AGNIFSgroup interprets the distinct flux distributions and kinematics ofmolecular gas and [Fe ii ] as an indication that molecular gas israther a tracer of the feeding of the AGN while [Fe ii ] is morelikely a tracer of the feedback (e.g., Ri ff el et al. 2006, 2008; Rif-fel & Storchi-Bergmann 2011a,b).Rotational and vibrational H emission lines are in partic-ular the dominant way to cool warm molecular gas. Line ratiosof these transitions and ionized hydrogen emission lines give im-portant information on their excitation mechanisms (Davies et al.2003, 2005). In the case of HE 1029-1831, they indicate a sig-nificant contribution of currently occuring star formation and / oryoung stellar populations in the central region.In general, the H emission lines can be excited by twokinds of processes: (1) thermal processes, e.g. produced by X-ray (Maloney et al. 1996) or by shock heating (Hollenbach &McKee 1989) and (2) non-thermal processes like UV fluores-cence (Black & van Dishoeck 1987): UV-photons with 912Å <λ < molecule in the Lyman- andWerner bands, exciting the next two electronic levels ( B Σ − X Σ and C Π − X Σ ). With a probability of 90%, a decay into a boundbut excited rovibrational level within the electronic ground level X Σ will take place. By this mechanism, H rovibrational levelswill be populated, which could not be populated by collisions.Possible sources are OB stars or strong AGN continuum emis-sion.Di ff erent excitation mechanisms can be discriminated by ob-serving H line ratios. In particular, the H line ratio 2-1S(1) / / line ratios 2-1S(1) / / Article number, page 9 of 16 & A proofs: manuscript no. he1029 more, we indicate the positions of models for thermal UV ex-citation (Sternberg & Dalgarno 1989), UV fluorescence (Black& van Dishoeck 1987), X-ray heating (Draine & Woods 1990),shock-heating (Brand et al. 1989), and a thermal emission curve.The position in the diagram shows that the H excitation in thetwo apertures cannot be explained by thermal excitation alone.A significant contribution from UV fluorescence is necessary.This is in good consistency with star formation activity in thetwo spots. Previous studies show that in the case of AGN, colli-sional excitation, e.g. by interaction of the radio-jet with the cir-cumnuclear interstellar medium, might be more important thanthe fluorescent process (Ri ff el et al. 2006, 2010). However, otherstudies see strong evidence that fluorescent excitation plays in-deed a major role in the H excitation in ULIRGs (Davies et al.2003) and AGNs (Davies et al. 2005), in concordance with ourfindings.The rotational temperature can be determined from two or-tho / para lines that belong to the same vibrational level, whereasthe vibrational temperature can be determined by connecting twotransitions with same J but from consecutive v levels: T rot( v = = . + ln (cid:16) f − f − (cid:17) (6) T vib = . + ln (cid:16) f − f − (cid:17) . (7)Using the flux measurements listed in Tab. 2, we estimate a ro-tational temperature of T rot = (1100 ± T rot = (900 ± T vib (spot1) = (3700 ± T vib (spot2) = (3200 ± molecules to be in local thermal equilibrium. In this case, T vib ≈ T rot . In particular, Rodríguez-Ardila et al. (2005) and Ri ff el et al.(2013) show that star forming galaxies show T vib (cid:38) T rot whileAGNs and LINERs tend to have similar values for T vib and T rot .In HE 1029-1831, the temperatures di ff er significantly from eachother. This points at a significant contribution from the non-thermal process of UV-fluorescence, consistent with star forma-tion in the circumnuclear region.Further support for this hypothesis is given by the line ra-tios of log([Fe ii ] / Br γ ) that are expected to be < . ii ] / Br γ ) nucl = .
04, log([Fe ii ] / Br γ ) spot1 = − .
01, and log([Fe ii ] / Br γ ) spot2 = . (1 − / Br γ )and log([Fe ii ] λ . µ m / Pa β ). The diagnostic is supposedto distinguish between excitation from pure photoioniza-tion and from pure shocks and shows a transition fromstarburst galaxies to LINERS, passing AGNs where bothexcitation mechanisms are of importance. Since we arelacking J -band information, we use the conversion factors[Fe ii ] λ . µ m / [Fe ii ] λ . µ m = .
744 (Nussbaumer &Storey 1988) and Pa α/ Pa β = .
05 (Osterbrock & Fer-land 2006). The line ratios [log(H / Br γ ) , log([Fe ii ] / Pa β )] = [ − . , − .
60] for the nuclear region, [ − . , − .
66] for spot1,and [ − . , − .
59] for spot2 are consistent with a mixture ofstarburst activity and AGN excitation (see Fig. 8 right). Previousstudies (e.g., Ri ff el et al. 2010; Schönell et al. 2014) show thatthe centers of active galaxies often show line ratios characteristicfor a mixture of starburst and AGN excitation. A su ffi cient reservoir of gas is necessary to fuel star formationand black hole accretion. We show that the LLQSO HE 1029-1831 is rich in ionized and molecular gas.We estimate the mass of ionized hydrogen as M H ii = m p n e V H ii with the electron density n e and the volume V H ii of theemitting region. Using the line coe ffi cients from Osterbrock &Ferland (2006), assuming an electron temperature of T = Kand a density in the range 10 < n e < cm − , we can calculatethe Pa α flux as f Pa α = (cid:33) j Pa α d Ω d V π d = π (cid:32) π j H β n (cid:33) (cid:32) j Pa α j H β (cid:33) n V H ii D (8) ≈ . × − n V H ii D erg cm − s − (9)where D is the distance to the galaxy in cm. The ionized gasmass is then given by M H ii ≈ . × (cid:32) f Pa α W m − (cid:33) (cid:32) D Mpc (cid:33) (cid:18) n e cm − (cid:19) − M (cid:12) (10)In an aperture with radius 4 (cid:48)(cid:48) , the flux in Pa α is f Pa α ( r ≤ (cid:48)(cid:48) ) = × − W m − . Assuming an electron density of n e =
100 cm − , this corresponds to a mass in ionized hydrogenof M H ii ( r ≤ (cid:48)(cid:48) ) = . × M (cid:12) .The mass of warm H can be estimated by M H ≈ . × (cid:32) F H W m − (cid:33) (cid:32) D Mpc (cid:33) M (cid:12) (11)following Scoville et al. (1982), Wolniewicz et al. (1998) andRi ff el et al. (2010). In an aperture of radius 4 (cid:48)(cid:48) , we measure aH (1-0)S(1) flux f H = × − W m − . This corresponds toa warm H gas mass of 4700 M (cid:12) . We see that the mass of ionizedgas is a factor of ≈ − (Ri ff el et al. 2014b, and referencestherein). With the conversion factor of Mazzalay et al. (2013b), M H (cold) / M H (warm) = (0 . − . × , we derive a cold gasmass of (1 . − . × M (cid:12) . Other cold-to-warm H gas massratios have been found by Dale et al. (2005) (10 − ) or MüllerSánchez et al. (2006) (2 . × with a 1 σ uncertainty of a factorof 2) which result in cold gas masses that are consistent with ourresults.Krips et al. (2007b) and Moser et al. (in prep.) derive themolecular gas mass from CO(1-0) measurements. They estimatemasses of a few to several 10 M (cid:12) , which is in good consistencywith our results.For NUGA sources, cold molecular gas masses have beencalculated from CO-emission, ranging from 2 × − × M (cid:12) with typical masses of the order of several 10 M (cid:12) (Moseret al. 2012, and references therein). The AGNIFS group ob-tained masses in a range of 66 M (cid:12) ≤ M H ≤ M (cid:12) and0 . × M (cid:12) ≤ M H ii ≤ × M (cid:12) in the nuclear regionsof nearby galaxies (Ri ff el et al. 2008, 2009; Storchi-Bergmannet al. 2009; Ri ff el et al. 2010; Schönell et al. 2014). We con-clude that the LLQSO HE 1029-1831 has a large reservoir ofcold molecular gas that is needed and apparently used for both,star formation and black hole fueling. Article number, page 10 of 16erold Busch et al.: LLQSO sample - Tracing circumnuclear star formation in HE 1029-1831 with SINFONI (2-1)S(1)/H (1-0)S(1)0.00.51.01.52.02.53.0 H ( - ) S ( ) / H ( - ) S ( ) Thermal non-Thermal3000KX-RaysShocks spot 1spot 2 T vib T r o t log(H / Br γ ) l og ( [ F e II ] / P a β ) SB AGN LINER
SBLINERSy1Sy2SN nuclearspot1spot2
Fig. 8.
Left:
Molecular hydrogen diagnostic diagram with 2-1S(1) / / Right:
Diagnostic diagram with line ratios 1-0S(1) λ . µ m / Br γ vs. [Fe ii ] λ . µ m / Pa β (for details see text). The positions of three apertures are indicated by circles. Open symbols correspond to literaturevalues from Larkin et al. (1998); Dale et al. (2004); Rodríguez-Ardila et al. (2004, 2005). The lines indicate regions that are typically populatedby starburst galaxies, AGNs, and LINERs resp. Dashed lines: Rodríguez-Ardila et al. (2005); Ri ff el et al. (2010), dotted lines: Ri ff el et al. (2013). In our previous study (Busch et al. 2014), we found that the ob-served LLQSOs do not follow published M BH − L bulge relationsfor inactive galaxies and speculated about the influence of youngor intermediate-age stellar populations in the bulge lowering themass-to-light ratio. Here, we estimate star formation rates andthe dynamical mass of the bulge. Furthermore, we use the equiv-alent width of Br γ , the supernova rate, and the mass-to-light ratioas diagnostics to contrain possible star formation histories andestimate the impact of star formation in the bulge. Since dust grains absorb the radiation of young stars and re-radiate it in the far-infrared (FIR), the FIR luminosity L FIR isa tracer of star formation activity on 100-300 Myr timescales.From the IRAS faint source catalog (Moshir et al. 1992),we calculate the FIR luminosity L FIR = π (cid:0) D L / Mpc (cid:1) × . × (2 . f + f ) W = . × W, where f λ is the fluxdensity at 60 µ m and 100 µ m resp. (Helou et al. 1988; Sanders &Mirabel 1996). With the calibration of Panuzzo et al. (2003), weestimate a global star formation rate ofSFR FIR = L FIR . × W = . M (cid:12) yr − (12)This FIR luminosity could be heavily a ff ected by the AGN emis-sion. Therefore, the star formation rate should only be taken asan upper limit.Furthermore, we calculate the FIR colors log( f / f ) = . ± .
13, log( f / f ) = . ± .
05, log( f / f ) = . ± .
12, and log( f / f ) = . ± .
12. According to Dopita et al.(1998); Kewley et al. (2000), these ratios are in the range of AGNwith more than 90% contribution to the FIR luminosity from starformation. Together with SFR
FIR calculated above, this impliesthat HE 1029-1831 is actively star forming.With our high-resolution data, we can calculate current starformation rates in the spatially resolved nuclear regions. We use Br γ as a star formation indicator according to Panuzzo et al.(2003)SFR M (cid:12) yr − = L (Br γ )1 . × W . (13)In the two o ff -nuclear spots, we measure SFR spot1 = . M (cid:12) yr − and SFR spot2 = . M (cid:12) yr − . This corresponds tostar formation rate densities of Σ SFR , spot1 = M (cid:12) yr − kpc − and Σ SFR , spot2 = M (cid:12) yr − kpc − . Our aperture has a radiusof 0 (cid:48)(cid:48) . =
120 pc. Thus, our values are in agreement withthe expected star formation rate densities in Seyfert galaxiesof (1 − M (cid:12) yr − kpc − on hundreds of parsecs scales and(50 − M (cid:12) yr − kpc − on tens of parsecs scales (Valencia-S.et al. 2012a, and references therein).Furthermore, we can compare the star formation rate den-sities to the gas mass densities (estimated from the hot molec-ular gas mass) and find log (cid:16) Σ gas , spot1 / M (cid:12) pc − (cid:17) = . (cid:16) Σ gas , spot2 / M (cid:12) pc − (cid:17) = .
6. This means that the star forma-tion regions follow the Schmidt law (Kennicutt 1998) and thehigh gas mass is e ffi ciently transformed into stars. The combination of the image analysis that provides us withthe e ff ective radius and the kinematical analysis that provides uswith the stellar velocity dispersion, allows us to calculate the dy-namical mass of the central bulge component. Sani et al. (2011),following Cappellari et al. (2006), adopt M dyn = κ r e σ G (14)with κ = ff ectiveradius r e . With the values for r e and σ e derived in Sections 3.1and 3.3, this results in M dyn = × M (cid:12) .From the H flux in an aperture with radius r e , we can es-timate the cold H gas mass in the bulge region (see Sect. 4.3). Article number, page 11 of 16 & A proofs: manuscript no. he1029
We multiply with 1.36 in order to allow for the contribution ofHelium gas and obtain (1 . − . × M (cid:12) . This correspondsto a gas fraction of 20% − M ∗ = (1 . − . × M (cid:12) .Several correlations between the black hole mass and the(stellar) bulge mass have been found in the last years. We usethe dynamical bulge mass estimate to find the correspondingblack hole mass: log( M BH / M (cid:12) ) = . − . M BH − L bulge relations. This can be interpreted ashint that the deviation of LLQSOs from the M BH − L bulge relationobserved in Busch et al. (2014) is - at least in the case of HE1029-1831 - primarily caused by the higher luminosity of youngstellar populations rather than a significantly undermassive blackhole. To follow this trace, we use S tarburst
99 (Leitherer et al. 1999;Vázquez & Leitherer 2005; Leitherer et al. 2010, 2014) to modelthe mass-to-light ratio ( M ∗ / L K ratio), Br γ equivalent width, andsupernova rate which can then be compared to observations(Davies et al. 2006, 2007; Levesque & Leitherer 2013). We showthat the bulge is dominated by young / intermediate-age popu-lations ( ∼
100 Myr) and find traces for a declining starburst ∼
200 Myr ago.For the simulations, we use the Padova tracks with AGB starsat solar metallicity and a Kroupa IMF. We consider five di ff erentstar formation histories: an instantaneous starburst, single star-bursts with decay times of τ SF =
10 Myr, 50 Myr, and 100 Myras well as continuous star formation. The decaying starburstswere simulated by adding up instantaneous starbursts scaled toexponentially decaying intensity with the respective time scales. Br γ equivalent width The equivalent width of the stellar Br γ emission is supposed to be a good estimator for the stellar agesince it decreases monotonically with time (see Fig. 9, left).However, the measurement could be a ff ected by the nonstellarAGN continuum. Davies et al. (2007) conclude that every stellarpopulation that contains late-type stars should show an equiva-lent width of the stellar absorption feature CO(2-0) at λ . µ mof W CO(2 − ≈
12 Å independent of the star formation historyand age with an uncertainty of about 20%. Thus, they suggestthat the fraction of the stellar light, contaminated by nonstellar-AGN continuum, can be estimated by f stellar = W obs / W intr withthe observed equivalent width of CO(2-0), W obs , and the intrinsic W intr =
12 Å.We use the definitions of Origlia et al. (1993) and measurethe absorption in the interval 2 . µ m − . µ m and the con-tinuum between 2 . µ m − . µ m. This results in equiv-alent widths of CO(2-0) of W CO , nucl = . W CO , spot1 = . W CO , spot2 = . f stellar , nucl = f stellar , spot1 = f stellar , spot2 = γ equivalent width in the three apertures. Aftercorrecting for the non-stellar contribution, we gain W Br γ, nucl , corr = (8 . ± .
2) Å, W Br γ, spot1 , corr = (15 . ± .
5) Å, and W Br γ, spot2 , corr = (9 . ± .
4) Å.
Supernova rate
The forbidden iron transition [Fe ii ] λ . µ mis a good shock tracer. Since we see no indication for jets andoutflows, we assume that the line is mainly excited by shocksfrom supernovae. The supernova rate can then be calculatedSNR Cal97 = . L [Fe ii ] W yr − (15)following Calzetti (1997) or, following Alonso-Herrero et al.(2003),SNR AH03 = . L [Fe ii ] W yr − . (16)We measure the [Fe ii ] emission in an aperture corresponding tothe half-light radius r e = (cid:48)(cid:48) .
59 and get L [Fe ii ] = . × W.This corresponds to supernova rates of SNR
Cal97 = .
055 yr − and SNR AH03 = .
083 yr − . In Fig. 9 (right), we show the super-nova rate divided by the stellar mass (normalized by 10 M (cid:12) ) assimulated by S tarburst
99. As bulge mass estimate, we use thedynamical mass derived above.To estimate the stellar mass in the two spots, we assumethat the bulge mass is smoothly distributed following a Sérsiclaw with the parameters derived in the bulge-disk decomposi-tion. The mass fraction of the spots can then easily be esti-mated whereas we determine the supernova rate from the [Fe ii ]flux in the apertures. The supernova rates in the two spots areSNR(spot1) = . − .
010 yr − and SNR(spot2) = . − .
008 yr − . The supernova rates should be taken as upper limitsonly, since contamination from the AGN cannot be excluded. Mass-to-light ratio
Figure 10 shows that the ratio of the stellarmass to the K -band luminosity can be used as a diagnostic for theage of the stellar population for ages ≥ yr − since it increasesmonotonically with time. However, two caveats have to be keptin mind: First, we cannot measure the stellar mass directly. In-stead, we will take the dynamical mass corrected for the coldgas mass as an estimate. Second, the model assumes that onlyone stellar component is discussed. However, we expect to havea mixture of an old population ( ≈ yr) that could be mixedwith one or more young populations that have been formed inlater starbursts. The age estimate from the mass-to-light ratio canthus only be an upper limit for the young population. With thebulge mass M bulge ≈ (1 − × M (cid:12) and the bulge magnitude M bulge = − .
57 derived from the bulge-disk decomposition inSect. 3.1, we derive a mass-to-light ratio of M ∗ / L K = . − . tarburst
99 models. From these diagrams, immedi-ate conclusions can be drawn: From the equivalent width of Br γ and the supernova rate, we can rule out an instantaneous star-burst scenario for spots 1 and 2 since from W Br γ , we wouldestimate a starburst age of about 5-6 Myr but in this case thesupernova rate should be much higher than observed. Continu-ous star formation is also rather inconsistent with our models:From W Br γ , we expect an age of around 800 Myr, while fromthe supernova rate the age would be around 200-500 Myr. Thefavored model is an exponentially decaying starburst with timescale τ SF = −
100 Myr that began 100-200 Myr ago.For the overall bulge, we can again exclude an instantaneousstarburst model from the supernova rates. The favored model forthe two spots is consistent with supernova rates and the mass-to-light ratio for the overall bulge, indicating that the bulge isdominated by the star formation regions. However, continuous
Article number, page 12 of 16erold Busch et al.: LLQSO sample - Tracing circumnuclear star formation in HE 1029-1831 with SINFONI age [yrs]110100500 W B r γ [] spot 1spot 2 instantaneous τ SF =10 Myr τ SF =50 Myr τ SF =100 Myrcontinuous age [yrs] -3 -2 -1 S N R [ y r − ] / M ∗ [ M fl , K ] bulgespot1 spot2 instantaneous τ SF =10 Myr τ SF =50 Myr τ SF =100 Myrcontinuous Fig. 9.
Results from the S tarburst
99 simulations. Br γ equivalent width ( left ) and supernova rate ( right ) as a function of age for five star formationhistories: Instantaneous starburst (“fixed mass” in S tarburst star formation that began around 500 Myr ago is also consistentwith the data.We conclude that we observe an intermediate-age stellarpopulation with an age of around 100 Myr in the two spotsthat dominate the overall bulge emission. The star formationtime scale is of the order of 50-100 Myr. This implies that thestar formation is declining, in accordance with the derived cur-rent star formation rates that were high compared to quiescentgalaxies but not as extraordinary as expected for a starburstgalaxy, and with the high value derived from the FIR luminos-ity that is sensitive to star formation on timescales of 100-300Myr. Furthermore, this is in good consistency with the findingof Davies et al. (2007) that AGNs accreting at lower e ffi ciency( λ ≤ .
1) have younger starbursts ( ≤ −
100 Myr) while AGNsaccreting at higher e ffi ciency ( λ ≥ .
1) have older starbursts( ≥ −
100 Myr).To conclude, the three diagnostics hint at star formation as animportant factor in the central region of the LLQSO HE 1029-1831. Although the star formation activity has already reachedits maximum more than 100 Myr ago, the e ff ect is still visible:The presence of intermediate-age stellar population lowers themass-to-light ratio by a factor of 10 compared to a bulge purelyconsisting of an old stellar population (see Fig. 10).Interestingly, the dynamical mass estimate that was derivedhere is consistent with predictions from the M BH − M bulge re-lations. Thus, at least in the particular case of HE 1029-1831,we conclude that the observed o ff set of LLQSOs from the M BH − L bulge relations for inactive galaxies is caused by recentstar formation activity in the central region.An alternative explanation is the presence of a pseudo-bulgewhich is supported by the low Sérsic index. Pseudo-bulges arethought not to evolve from galaxy mergers but from secular evo-lution, therefore not obeying the relations of classical bulgesand ellipticals (Kormendy & Kennicutt 2004, and referencestherein). Other authors suggest that there exist two relations fortwo di ff erent formation mechanisms (gaseous formation mecha-nism and dry mergers, Graham & Scott 2013), meaning that thereplacement from the relation is not necessarily caused by un-dermassive black holes or overluminous bulges but by di ff erentformation processes. In particular, the M BH − L bulge relation forgalaxies which originate from gaseous formation mechanismshas a large scatter. Within this scatter, our galaxy might follow age [yrs]0.010.11 M ∗ / L K [ M fl / L fl , K ] bulgeinstantaneous τ SF =10 Myr τ SF =50 Myr τ SF =100 Myrcontinuous Fig. 10. K -band mass-to-light ratio as a function of the age of the stellarpopulation, result from the S tarburst
99 simulations. The M ∗ / L K ratiorange derived from the dynamical mass and the bulge-disk decomposi-tion (see Sect. 4.4) is indicated in gray. the relation. However, in our previous study (Busch et al. 2014)with 11 objects, we show that LLQSOs (like HE 1029-1831)are systematically shifted from the Graham&Scott-relation to theright.We emphasize that both explanations are not inconsistentwith our findings since in both scenarios the central compo-nent (“pseudo-bulge” or “Sérsic spheroid”) is expected to be richin gas and to contain younger stellar populations than classical(dry) bulges.
5. Summary and conclusions
We have observed the inner 6 (cid:48)(cid:48) . (cid:48)(cid:48) . (cid:48)(cid:48) . =
110 pc with AO, using the integral-field spectrographSINFONI mounted at the VLT UT4 in Paranal. The main resultsof our analysis are the following:
Article number, page 13 of 16 & A proofs: manuscript no. he1029 – The H ii emission shows a di ff erent flux distribution than thestellar continuum. Two gas spiral arms are located within thestellar bar. At smaller scales, a patchy circumnuclear ring isseen. – The flux distribution of the shock tracer [Fe ii ] follows theH ii insofar as the emission is strongest in the points wherethe gas spiral arms meet the ring. Diagnostic line ratios ase.g. log(H / Br γ ), log([Fe ii ] / Br γ ), or molecular hydrogenline ratios are indicative for star formation regions. – The mass of ionized hydrogen in the inner 3 . M H ii = . × M (cid:12) . The mass of cold molecular hydrogen within1 . M H , cold = (1 . − . × M (cid:12) . This is in goodagreement with H -mass estimates from CO measurementsand comparable to gas masses derived for NUGA galaxies.We conclude that HE 1029-1831 has a massive reservoir ofgas that is available for star formation and AGN fueling. – A bulge-disk decomposition with B udda reveals that thebulge luminosity does not follow published M BH − L bulge re-lations for inactive galaxies. This is expected for LLQSOs(Busch et al. 2014). The dynamical mass of the bulge, how-ever, is in good agreement with common M BH − M bulge rela-tions. – From the M ∗ / L K ratio and diagnostics like the Br γ equiva-lent width and the supernova rates, we find evidence for thepresence of young / intermediate-age stellar populations in thecircumnuclear region. The favored model is an exponentiallydecaying starburst with time scale 50-100 Myr that beganaround 100-200 Myr ago. The light of the bulge componentis dominated by this intermediate-age stellar population.We conclude that circumnuclear star formation is a crucialfactor in HE 1029-1831 that has to be taken into account forany thorough analysis (see also, e.g., Jahnke et al. 2004; CidFernandes et al. 2004). We find evidence that, at least in the caseof HE 1029-1831, the o ff set from the M BH − L bulge relations couldbe explained by an overluminosity of the bulge component dueto stellar populations that are younger than commonly observedin bulges.As pointed out in Sect. 4.1, the Eddington ratio ( λ (cid:38) .
06) ishigh, meaning that the black hole is in a growing phase. Though,since the black hole mass and - at least the dynamical - massobey published M BH − M bulge relations, from our data it seemsrather unlikely that the o ff set from the M BH − L bulge relation iscaused by a significantly undermassive black hole that wouldgrow “towards” the M BH − L bulge relation.With a larger sample of LLQSOs that have high-resolutiondata, in spatial and spectral dimension, we will be able to furtherinvestigate the star-formation properties in the circumnuclear re-gions. Acknowledgements.
The authors thank the anonymous referee for the helpful re-port. G. Busch thanks Marcus Bremer for many fruitful discussions. This workwas supported in part by the Deutsche Forschungsgemeinschaft (DFG) via SFB956. G. Busch is member of the Bonn-Cologne Graduate School of Physicsand Astronomy (BCGS). S. Smaji´c is member of the International Max PlanckResearch School (IMPRS) for Astronomy and Astrophysics Bonn / Cologne.J. Scharwächter acknowledges the European Research Council for the AdvancedGrant Program Number 267399-Momentum. M. Valencia-S. received fundingfrom the European Union Seventh Framework Programme (FP7 / References
Alonso-Herrero, A., Rieke, G. H., Rieke, M. J., & Kelly, D. M. 2003, AJ, 125,1210Alonso-Herrero, A., Rieke, M. J., Rieke, G. H., & Ruiz, M. 1997, ApJ, 482, 747 Antonucci, R. 1993, ARA&A, 31, 473Bennert, V. N., Auger, M. W., Treu, T., Woo, J.-H., & Malkan, M. A. 2011, ApJ,726, 59Bentz, M. C., Peterson, B. M., Netzer, H., Pogge, R. W., & Vestergaard, M. 2009,ApJ, 697, 160Bertram, T., Eckart, A., Fischer, S., et al. 2007, A&A, 470, 571Black, J. H. & van Dishoeck, E. F. 1987, ApJ, 322, 412Böker, T., Falcón-Barroso, J., Schinnerer, E., Knapen, J. H., & Ryder, S. 2008,AJ, 135, 479Bonnet, H., Abuter, R., Baker, A., et al. 2004, The Messenger, 117, 17Brand, P. W. J. L., Toner, M. P., Geballe, T. R., et al. 1989, MNRAS, 236, 929Busch, G., Zuther, J., Valencia-S., M., Moser, L., & Eckart, A. 2012, in Pro-ceedings of Nuclei of Seyfert galaxies and QSOs - Central engine & con-ditions of star formation (Seyfert 2012). 6-8 November, 2012. Online athttp: // pos.sissa.it / cgi-bin / reader / conf.cgi?confid = / Infrared Ground-based Telescopes, ed.M. Iye & A. F. M. Moorwood, 1548–1561Falcón-Barroso, J., Ramos Almeida, C., Böker, T., et al. 2014, MNRAS, 438,329Ferrarese, L. & Merritt, D. 2000, ApJ, 539, L9Fischer, S., Iserlohe, C., Zuther, J., et al. 2006, A&A, 452, 827Freeman, K. C. 1970, ApJ, 160, 811Gadotti, D. A. 2008, MNRAS, 384, 420Gadotti, D. A. 2009, MNRAS, 393, 1531García-Burillo, S., Combes, F., Hunt, L. K., et al. 2003, A&A, 407, 485García-Burillo, S., Combes, F., Schinnerer, E., Boone, F., & Hunt, L. K. 2005,A&A, 441, 1011Graham, A. W. & Driver, S. P. 2007, ApJ, 655, 77Graham, A. W., Onken, C. A., Athanassoula, E., & Combes, F. 2011, MNRAS,412, 2211Graham, A. W. & Scott, N. 2013, ApJ, 764, 151Gültekin, K., Richstone, D. O., Gebhardt, K., et al. 2009, ApJ, 698, 198Häring, N. & Rix, H.-W. 2004, ApJ, 604, L89Helou, G., Khan, I. R., Malek, L., & Boehmer, L. 1988, ApJS, 68, 151Hollenbach, D. & McKee, C. F. 1989, ApJ, 342, 306Hopkins, P. F., Richards, G. T., & Hernquist, L. 2007, ApJ, 654, 731Husemann, B., Jahnke, K., Sánchez, S. F., et al. 2014, MNRAS, 443, 755Husemann, B., Wisotzki, L., Sánchez, S. F., & Jahnke, K. 2013, A&A, 549, A43Jahnke, K., Kuhlbrodt, B., & Wisotzki, L. 2004, MNRAS, 352, 399Kaspi, S., Smith, P. S., Netzer, H., et al. 2000, ApJ, 533, 631Kennicutt, Jr., R. C. 1998, ApJ, 498, 541Kewley, L. J., Heisler, C. A., Dopita, M. A., & Lumsden, S. 2001, ApJS, 132, 37Kewley, L. J., Heisler, C. A., Dopita, M. A., et al. 2000, ApJ, 530, 704Kim, D., Im, M., & Kim, M. 2010, ApJ, 724, 386Kim, M., Ho, L. C., Peng, C. Y., et al. 2008, ApJ, 687, 767König, S., Eckart, A., García-Marín, M., & Huchtmeier, W. K. 2009, A&A, 507,757Kormendy, J., Bender, R., & Cornell, M. E. 2011, Nature, 469, 374Kormendy, J. & Ho, L. C. 2013, ARA&A, 51, 511Kormendy, J. & Kennicutt, Jr., R. C. 2004, ARA&A, 42, 603Krips, M., Eckart, A., Krichbaum, T. P., et al. 2007a, A&A, 464, 553Krips, M., Eckart, A., Neri, R., et al. 2007b, A&A, 464, 187Lançon, A., Hauschildt, P. H., Ladjal, D., & Mouhcine, M. 2007, A&A, 468, 205Larkin, J. E., Armus, L., Knop, R. A., Soifer, B. T., & Matthews, K. 1998, ApJS,114, 59Läsker, R., Ferrarese, L., van de Ven, G., & Shankar, F. 2014, ApJ, 780, 70
Article number, page 14 of 16erold Busch et al.: LLQSO sample - Tracing circumnuclear star formation in HE 1029-1831 with SINFONI
Leitherer, C., Ekström, S., Meynet, G., et al. 2014, ApJS, 212, 14Leitherer, C., Ortiz Otálvaro, P. A., Bresolin, F., et al. 2010, ApJS, 189, 309Leitherer, C., Schaerer, D., Goldader, J. D., et al. 1999, ApJS, 123, 3Levesque, E. M. & Leitherer, C. 2013, ApJ, 779, 170Magorrian, J., Tremaine, S., Richstone, D., et al. 1998, AJ, 115, 2285Mahony, E. K., Croom, S. M., Boyle, B. J., et al. 2010, MNRAS, 401, 1151Maiolino, R., Rieke, G. H., & Rieke, M. J. 1996, AJ, 111, 537Maloney, P. R., Hollenbach, D. J., & Tielens, A. G. G. M. 1996, ApJ, 466, 561Marconi, A. & Hunt, L. K. 2003, ApJ, 589, L21Markwardt, C. B. 2009, in Astronomical Society of the Pacific Conference Se-ries, Vol. 411, Astronomical Data Analysis Software and Systems XVIII, ed.D. A. Bohlender, D. Durand, & P. Dowler, 251Mazzalay, X., Rodríguez-Ardila, A., Komossa, S., & McGregor, P. J. 2013a,MNRAS, 430, 2411Mazzalay, X., Saglia, R. P., Erwin, P., et al. 2013b, MNRAS, 428, 2389Moorwood, A., Cuby, J.-G., Biereichel, P., et al. 1998, The Messenger, 94, 7Moser, L., Zuther, J., Busch, G., Valencia-S., M., & Eckart, A. 2012, in Proceed-ings of Nuclei of Seyfert galaxies and QSOs - Central engine & conditions ofstar formation (Seyfert 2012). 6-8 November, 2012. Max-Planck-Insitut fürRadioastronomie (MPIfR), Bonn, Germany. Online at http: // pos.sissa.it / cgi-bin / reader / conf.cgi?confid = ff el, R., Rodríguez-Ardila, A., Aleman, I., et al. 2013, MNRAS, 430, 2002Ri ff el, R. A., Ho, L. C., Mason, R., et al. 2014a, ArXiv e-printsRi ff el, R. A. & Storchi-Bergmann, T. 2011a, MNRAS, 411, 469Ri ff el, R. A. & Storchi-Bergmann, T. 2011b, MNRAS, 417, 2752Ri ff el, R. A., Storchi-Bergmann, T., Dors, O. L., & Winge, C. 2009, MNRAS,393, 783Ri ff el, R. A., Storchi-Bergmann, T., & Nagar, N. M. 2010, MNRAS, 404, 166Ri ff el, R. A., Storchi-Bergmann, T., Winge, C., & Barbosa, F. K. B. 2006, MN-RAS, 373, 2Ri ff el, R. A., Storchi-Bergmann, T., Winge, C., et al. 2008, MNRAS, 385, 1129Ri ff el, R. A., Vale, T. B., Storchi-Bergmann, T., & McGregor, P. J. 2014b, MN-RAS, 442, 656Rodríguez-Ardila, A., Pastoriza, M. G., & Donzelli, C. J. 2000, ApJS, 126, 63Rodríguez-Ardila, A., Pastoriza, M. G., Viegas, S., Sigut, T. A. A., & Pradhan,A. K. 2004, A&A, 425, 457Rodríguez-Ardila, A., Ri ff el, R., & Pastoriza, M. G. 2005, MNRAS, 364, 1041Sanders, D. B. & Mirabel, I. F. 1996, ARA&A, 34, 749Sani, E., Marconi, A., Hunt, L. K., & Risaliti, G. 2011, MNRAS, 413, 1479Savorgnan, G., Graham, A. W., Marconi, A., et al. 2013, MNRAS, 434, 387Schönell, A. J., Ri ff el, R. A., Storchi-Bergmann, T., & Winge, C. 2014, MNRAS,445, 414Schulze, A. & Wisotzki, L. 2010, A&A, 516, A87Schulze, A., Wisotzki, L., & Husemann, B. 2009, A&A, 507, 781Scott, N., Graham, A. W., & Schombert, J. 2013, ApJ, 768, 76Scoville, N. Z., Hall, D. N. B., Ridgway, S. T., & Kleinmann, S. G. 1982, ApJ,253, 136Sersic, J. L. 1968, Atlas de galaxias australes, ed. Sersic, J. L.Skrutskie, M. F., Cutri, R. M., Stiening, R., et al. 2006, AJ, 131, 1163Smaji´c, S., Moser, L., Eckart, A., et al. 2014, A&A, 567, A119Sternberg, A. & Dalgarno, A. 1989, ApJ, 338, 197Storchi-Bergmann, T., McGregor, P. J., Ri ff el, R. A., et al. 2009, MNRAS, 394,1148Urry, C. M. & Padovani, P. 1995, PASP, 107, 803Valencia-S., M., Zuther, J., Eckart, A., et al. 2012a, A&A, 544, A129Valencia-S., M., Zuther, J., Eckart, A., et al. 2012b, in Proceedings of Nu-clei of Seyfert galaxies and QSOs - Central engine & conditions of starformation (Seyfert 2012). 6-8 November, 2012. Max-Planck-Insitut für Ra-dioastronomie (MPIfR), Bonn, Germany. Online at http: // pos.sissa.it / cgi-bin / reader / conf.cgi?confid = ff el, R. A., & Storchi-Bergmann, T. 2009, ApJS, 185, 186Wisotzki, L., Christlieb, N., Bade, N., et al. 2000, A&A, 358, 77Wolniewicz, L., Simbotin, I., & Dalgarno, A. 1998, ApJS, 115, 293Woo, J.-H., Schulze, A., Park, D., et al. 2013, ApJ, 772, 49Woo, J.-H. & Urry, C. M. 2002, ApJ, 579, 530Zuther, J., Iserlohe, C., Pott, J.-U., et al. 2007, A&A, 466, 451 Appendix A: Appendix
Article number, page 15 of 16 & A proofs: manuscript no. he1029 ∆ RA [arcsec]432101234 ∆ D E C [ a r c s e c ] Br γ - Flux F l u x [ − W m − ] ∆ RA [arcsec]432101234 ∆ D E C [ a r c s e c ] Br γ - Equivalent width E W [ ] ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] F l u x [ − W m − ] ∆ RA [arcsec]1.51.00.50.00.51.01.5 ∆ D E C [ a r c s e c ] E W [ ] Fig. A.1.
The same as Fig. 5 but for Br γγ