A Close Look at Star Formation around Active Galactic Nuclei
R. I. Davies, F. Mueller Sánchez, R. Genzel, L.J. Tacconi, E.K.S. Hicks, S. Friedrich,
Max Planck Institut für extraterrestrische Physik, Postfach 1312, 85741, Garching, Germany
A. Sternberg
School of Physics and Astronomy, Tel Aviv University, Tel Aviv 69978, Israel
ABSTRACT
We analyse star formation in the nuclei of 9 Seyfert galaxies at spatial resolutions down
to 0.085′′, corresponding to length scales of order 10 pc in most objects. Our data were taken
mostly with the near infrared adaptive optics integral field spectrograph SINFONI. The stellar
light profiles typically have size scales of a few tens of parsecs. In two cases there is unambiguous
kinematic evidence for stellar disks on these scales. In the nuclear regions there appear to have
been recent – but no longer active – starbursts in the last 10-300Myr. The stellar luminosity
is less than a few percent of the AGN in the central 10 pc, whereas on kiloparsec scales the
luminosities are comparable. The surface stellar luminosity density follows a similar trend in all
the objects, increasing steadily at smaller radii up to ∼ 1013 L⊙ kpc
−2 in the central few parsecs,
where the mass surface density exceeds 104 M⊙ pc
−2. The intense starbursts were probably
Eddington limited and hence inevitably short-lived, implying that the starbursts occur in multiple
short bursts. The data hint at a delay of 50–100Myr between the onset of star formation and
subsequent fuelling of the black hole. We discuss whether this may be a consequence of the role
that stellar ejecta could play in fuelling the black hole. While a significant mass is ejected by OB
winds and supernovae, their high velocity means that very little of it can be accreted. On the
other hand winds from AGB stars ultimately dominate the total mass loss, and they can also be
accreted very efficiently because of their slow speeds.
Subject headings: galaxies: active — galaxies: nuclei — galaxies: Seyfert — galaxies: starburst
— infrared: galaxies
1. Introduction
During recent years there has been increasing evidence for a connection between active galactic nuclei
(AGN) and star formation in the vicinity of the central black holes. This subject forms the central topic of
this paper, and is discussed in Sections 4 and 5.
A large number of studies have addressed the issue of star formation around AGN. Those which have
probed closest to the nucleus, typically on scales of a few hundred parsecs, have tended to focus on Seyferts
– notably Seyfert 2 galaxies – since these are the closest examples (Sarzi et al. 2007; Asari et el. 2007;
1Based on observations at the European Southern Observatory VLT (60.A-9235, 070.B-0649, 070.B-0664, 074.B-9012, 076.B-
0098).
http://arxiv.org/abs/0704.1374v2
– 2 –
González Delgado, et al. 2005; Cid Fernandes et al. 2004; González Delgado, et al. 2001; Gu et al. 2001;
Joguet et al. 2001; Storchi-Bergmann et al. 2001; Ivanov et al. 2000). The overall conclusion of these studies
is that in 30–50% cases the AGN is associated with young (i.e. age less than a few 100 Myr) star formation.
While this certainly implies a link, it does not necessarily imply any causal link between the two phenomena.
Instead, it could more simply be a natural consequence of the fact that both AGN and starburst require gas
to fuel them. And that in some galaxies this gas has fallen towards the nucleus, either due to an interaction
or secular evolution such as bar driven inflow.
One aspect which must be borne in mind when interpreting such results, and which has been pointed
out by Knapen (2004), is the discrepancy in the scales involved. AGN and starburst phemonena occur on
different temporal and spatial scales; and observations are sensitive to scales that are different again. For
example, star formation has typically been studied on scales of several kiloparsecs down to a few hundred
parsecs. In contrast, accretion of gas onto an AGN will occur on scales much less than 1 pc. Similarly, the
shortest star formation timescales that most observations are sensitive to are of order 100 Myr to 1 Gyr. On
the other hand, in this paper we show that the active phase of star formation close around a black hole is
typically rather less than 100 Myr. Correspondingly short accretion timescales for black holes are reflected
in the ages of jets which, for a sample of radio galaxies measured by Machalski et al. (2007), span a range
from a few to 100 Myr. In Seyfert galaxies the timescales are even shorter, as typified by NGC 1068 for which
Capetti et al. (1999) estimate the age of the jets to be only ∼ 0.1 Myr. That the putative causal connection
between AGN and starbursts might occur on relatively small spatial scales and short timescales can help us
to understand why no correlation has been found between AGN and (circum-)nuclear starbursts in general.
It is simply that the circumnuclear activity on scales greater than a few hundred parsecs is, in most cases,
too far from the AGN to influence it, or be strongly influenced by it (cf Heckman et al. 1997).
In this paper we redress this imbalance. While the optical spectroscopy pursued by many authors
allows a detailed fitting of templates and models to the stellar features, we also make use of established
star formation diagnostics and interpret them using starburst population synthesis models. Observing at
near infrared wavelengths has brought two important advantages. The optical depth is 10 times less than at
optical wavelengths, and thus our data are less prone to the effects of extinction which can be significant in
AGN. And we have employed adaptive optics to reach spatial resolutions of 0.1–0.2′′, bringing us closer to the
nucleus. Applying these techniques, we have already analysed the properties of the nuclear star formation in
a few objects (Davies et al. 2004a,b, 2006; Mueller Sánchez et al. 2006). Here we bring those data together
with new data on 5 additional objects. Our sample enables us to probe star formation in AGN from radii
of 1 kpc down to less than 10 pc. Our aim is to ascertain whether there is evidence for star formation on
the smallest scales we can reach; and if so, to constrain its star formation history. Ultimately, we look at
whether there are indications that the nuclear starburst and AGN are mutually influencing each other.
In §2 we describe the sample selection, observations, data reduction, PSF estimation, and extraction of
the emission and absorption line morphologies and kinematics. In §3 we discuss the observational diagnostics
and modelling tools. Brief analyses of the relevant facets of our new data for the individual objects are
provided in Appendix A, where we also summarise results of our previously published data, re-assessing
them where necessary to ensure that all objects are analysed in a consistent manner. The primary aims of
our paper are addressed in §4 and §5. In §4 we discuss global results concerning the existence and recent
history of nuclear star formation for our whole sample. In §5 we discuss the implications of nuclear starbursts
on the starburst-AGN connection. Finally, we present our conclusions in §6.
– 3 –
2. Sample, Observations, Data Processing
2.1. Sample Selection
The AGN discussed in this paper form a rather heterogeneous group. They include type 1 and type 2
Seyferts, ULIRGs, and even a QSO, and do not constitute a complete sample. In order to maximise the size
of the sample, we have combined objects on which we have already published adaptive optics near infrared
spectra with new observations of additional targets.
Source selection was driven largely by technical considerations for the adaptive optics (AO) system,
namely having a nucleus bright and compact enough to allow a good AO correction. This is actually a
strength since it means that 7 of the 9 AGN are in fact broad line objects – as given either by the standard
type 1 classification or because there is clear broad (FWHM > 1000km s−1) Brγ emission in our spectra.
Fig. 1 shows broad Brγ in K-band spectra of 3 AGN that are not usually classified as broad line galaxies.
This is in contrast to most other samples of AGN for which star formation has been studied in detail, and
avoids any bias that might arise from selecting only type 2 Syeferts. That there may be a bias arises from the
increasing evidence that the obscuration in perhaps half of type 2 AGN lies at kpc scales rather than in the
nucleus, which may be caused by spatially extended star formation in the galaxy disk (Brand et al. 2007;
Martinez-Sansigre et al. 2006; Rigby et al. 2006). Such AGN do not fit easily into the standard unification
scheme (and perhaps should not really be considered type 2 objects). Because broad lines can be seen in
the infrared, we know that we are seeing down to the nuclear region and hence our results are not subject
to any effects that this might otherwise introduce.
It is exactly broad line AGN for which little is known about the nuclear star formation, because the glare
of the AGN swamps any surrounding stellar light in the central arcsec. As a result, most studies addressing
star formation close to AGN have focussed on type 2 Seyferts. Adaptive optics makes it possible to confine
much of the AGN’s light into a very compact region, and to resolve the stellar continuum around it. The
use of adaptive optics does give rise to one difficulty when attempting to quantify the results in a uniform
way, due to the different resolutions achieved – which is a combination of both the distance to each object
(i.e. target selection) and the AO performance. As a result, the standard deviation around the logarithmic
mean resolution of our sample (excluding NGC 2992, see Section A) of 22 pc is a factor of 3. However, this
has enabled us to study the centers of AGN across nearly 3 orders of magnitude in spatial scale, from 1 kpc
in the more distant objects to only a few parsecs in the nearby objects with the best AO correction.
2.2. Observations & Reduction
A summary of our observations is given in Table 1. A description of observations and processing of the
new data is given below.
Data for IRAS 05189-2524 and NGC 1068 were taken in December 2002 at the VLT with NACO, an
adaptive optics near infrared camera and long slit spectrograph (Lenzen et al. 2003; Rousset et al. 2003).
Since IRAS 05189-2524 is nearly face on, there is no strongly preferred axis and the slit was oriented north-
south; for NGC 1068 two orientations were used, north-south and east-west. In all cases the slit width was
0.086′′, yielding a nominal resolution of R ∼ 1500 with the wide-field camera (pixel scale 0.054′′) and medium
resolution grism. The galaxy was nodded back and forth along the slit by 10′′ to allow sky subtraction. For
IRAS 05189-2524, 12 integrations of 300 sec were made; for NGC 1068 12 integrations of 300 sec were made
at one position angle, and 14 frames of 200 sec at the other. All data were reduced and combined, using
– 4 –
standard longslit techniques in IRAF, to make the final H-band spectrum.
Data for NGC 7469, NGC 2992, NGC 1097, NGC 1068, and NGC 3783 were taken during 2004–2005 at
the VLT with SINFONI, an adaptive optics near infrared integral field spectrograph (Eisenhauer et al. 2003;
Bonnet et al. 2004). Data were taken with various gratings covering the H and K bands either separately
(R ∼ 4000) or together (R ∼ 1500). The pixel scales were 0.125′′×0.25′′ or 0.05′′×0.1′′, depending on the
trade-offs between field of view, spatial resolution, and signal-to-noise ratio. Individual exposure times are in
the range 50–300 sec depending on the object brightness. Object frames were interspersed with sky frames,
usually using the sequence O-S-O-O-S-O, to facilitate background subtraction. The data were processed
using the dedicated spred software package (Abuter et al. 2006), which provides similar processing to that
for longslit data but with the added ability to reconstruct the datacube. The data processing steps are
as follows. The object frames are pre-processed by subtracting sky frames, flatfielding, and correcting bad
pixels (which are identified from dark frames and the flatfield). The wavemap is generated, and edges and
curvature of the slitlets are traced, all from the arclamp frame. The arclamp frame is then reconstructed
into a cube, which is checked to ensure that the calibration is good. The pre-processed object frames are
then also reconstructed into cubes, spatially shifted to align them using the bright nucleus as a reference,
and combined. In some cases the final cube was spatially smoothed using a 3 × 3 median filter. Estimation
of the spatial resolution (see below) was always performed after this stage.
In some cases, the strong near-infrared OH lines did not subtract well. With longer exposure times this
is to be expected since the timescale for variation of the OH is only 1–2 mins. If visual inspection of the
reconstructed cubes showed signs of over- or under-subtraction of the OH lines, these cubes were reprocessed
using the method described in Davies (2007a).
Standard star frames are similarly reconstructed into cubes. Telluric correction and flux calibration were
performed using B stars (K-band) or G2V stars (H-band). In addition, flux calibration was cross-checked in
3′′ apertures using 2MASS data, and in smaller 1–3′′ apertures using broad-band imaging from NACO or
HST NICMOS. Agreement between cubes with different pixel scales, and also with the external data, was
consistent to typically 20%.
2.3. PSF Estimation
There are a multitude of ways to derive the point spread function (PSF) from adaptive optics data,
five of which are described in Davies (2007b). With AGN, it is usually possible to estimate the PSF from
the science data itself, removing any uncertainty about spatial and temporal variations of the PSF due to
atmospheric effects. Typically one or both of the following methods are employed on the new data presented
here. If a broad emission line is detected, this will always yield a measure of the PSF since the BLR of
Seyfert galaxies has a diameter that can be measured in light days. Alternatively, the non-stellar continuum
will provide a sufficiently good approximation in all but the nearest AGN since at near infrared wavelengths
it is expected to originate from a region no more than 1–2 pc across.
In every case we have fit an analytical function to the PSF. Since the Strehl ratio achieved is relatively
low, even a Gaussian is a good representation. We have used a Moffat function, which achieves a better fit
because it also matches the rather broad wings that are a characteristic of partial adaptive optics correction.
The PSF measured for NGC 3227, which is shown in Fig. 1 of Davies et al. (2006), can be considered typical.
If one applies the concept of ‘core plus halo’ to this PSF, then the Gaussian fit would represent just the core
while the Moffat fit the entire ‘core plus halo’. Integrating both of these functions indicates that about 75%
– 5 –
of the flux is within the ‘core’, and it is thus this component which dominates the PSF. In this paper, a more
exact representation of the PSF is not needed since we have not performed a detailed kinematic analysis,
and we have simply used the Moffat to derive a FWHM for the spatial resolution. The resolutions achieved
are listed in Table 1.
2.4. Emission/Absorption Line Characterisation
The 2D distribution of emission and absorption features has been found by fitting a function to the
continuum-subtracted spectral profile at each spatial position in the datacube. The function was a convolu-
tion of a Gaussian with a spectrally unresolved template profile – in the case of emission lines it was an OH
sky emission line, and for stellar absorption features we made use of template stars observed in the same con-
figuration (pixel scale and grism). A minimisation was performed in which the parameters of the Gaussian
were adjusted until the convolved profile best matched the data. During the minimisation, pixels in the data
that consistently deviated strongly from the data were rejected. The uncertainties were boot-strapped using
Monte Carlo techniques, assuming that the noise is uncorrelated and the intrinsic profile is well represented
by a Gaussian. The method involves adding a Gaussian with the derived properties to a spectral segment
that exhibits the same noise statistics as the data, and refitting the result to yield a new set of Gaussian
parameters. After repeating this 100 times, the standard deviation of the center and dispersion were used
as the uncertainites for the velocity and line width.
The kinematics were further processed using kinemetry (Krajnović et al. 2006). This is a parameteri-
sation (i.e. a mathematical rather than a physical model) of the 2D field. As such, beam smearing is not
a relevant issue to kinemetry, which yields an analytical expression for the observed data. Of course, when
the coefficients of this expression are interpreted or used to constrain a physical model, then beam smearing
should be considered. Mathematically, the kinemetry procedure fits the data with a linear sum of sines and
cosines with various angular scalings around ellipses at each radius. We have used it for 3 purposes: to de-
termine the best position angle and axis ratio for the velocity field, to remove high order noise from the raw
kinematic extraction, and to recover the velocity and dispersion radial profiles. In all of the cases considered
here, the kinematic centre of the velocity field was assumed to be coincident with the peak of the non-stellar
continuum. In addition, the uniformity of the velocity field permitted us to make the simplifying assumption
of a single position angle and axis ratio – i.e. there is no evidence for warps or twisted velocity contours.
We then derived the position angle and inclination of the disk by minimising the A1 and B3 parameters
respectively (see Krajnović et al. 2006 for a description of these). The rotation curves were recovered by
correcting the measured velocity profile for inclination. We have assumed throughout the paper that the
dispersion is isotropic, and hence no inclination correction was applied to the dispersion that was measured.
The innermost parts of the kinematics derived as above are of course still affected by beam smearing.
In general, the central dispersion cannot necessarily be taken at face value since it may either be artificially
increased by any component of rotation included within the beam size, or decreased if neighbouring regions
within the beam have a lower dispersion. In the galaxies we have studied, there are two aspects which
mitigate this uncertainty: the rotation speed in the central region is much less than the dispersion and so
will not significantly alter it; and when estimating the central value we consider the trend of the dispersion
from large radii, where the effect of the beam is small, to the center. For the basic analyses performed here,
we have therefore adopted the central dispersion at face value. More detailed physical models for the nuclear
disks, which properly account for the effects of beam smearing, will be presented in future publications.
Lastly, we emphasize that the impact of the finite beam size on the derived rotation curve does not affect
– 6 –
our measurement of the dynamical mass. The reason is that, for all the dynamical mass estimates we make,
the mass is estimated at a radius much large than the FWHM of the PSF – as can be seen in the relevant
figures.
3. Quantifying the Star Formation
In this section we describe the tools of the trade used to analyse the data, and which lead us to the
global results presented in Section 4. Specific details and analyses for individual objects can be found in
Appendix A. We use the same methods and tools for all the objects to ensure that all the data are analysed
in a consistent manner.
Perhaps the most important issue is how to isolate the stellar continuum, which is itself a powerful
diagnostic. In addition, we use three standard and independent diagnostics to quantify the star formation
history and intensity in the nuclei of these AGN. These are the Brγ equivalent width, supernova rate, and
mass-to-light ratio. Much of the discussion concerns how we take into account the contribution of the AGN
when quantifying these parameters. We also consider what impact an incorrect compensation could have on
interpretation of the diagnostics.
We model these observational diagnostics using the stellar population and spectral synthesis code STARS
(e.g. Sternberg 1998; Sternberg et al. 2003; Förster Schreiber et al. 2003; Davies et al. 2003, 2005). This code
calculates the distribution of stars in the Hertzsprung-Russell diagram as a function of age for an assumed
star formation history. We usually assume an exponentially decaying star formation rate, which has an
associated timescale τSF. Spectral properties of the cluster are then computed given the stellar population
present at any time. We note that the model output of STARS is quantitatively similar to that from
version 5.1 of Starburst99, which unlike earlier versions does include AGB tracks (Leitherer et al. 1999;
Vazquez & Leitherer 2005). As discussed in detail below, particular predictions of STARS include, for ages
greater than 10 Myr: equivalent widths of WCO2−0 ∼ 12Å and WCO6−3 ∼ 4.5Å, and H-K color of 0.15 mag.
The equivalent quantities predicted by Starburst99 v5.1 are WCO2−0 ∼ 11Å and WCO6−3 ∼ 5Å, and H-K
color of 0.2 mag.
3.1. Isolating the stellar continuum
For small observational apertures a significant fraction of the K-band (and even H-band) continuum
can be associated with non-stellar AGN continuum. The AGN contribution can be estimated from a simple
measurement of the equivalent width of a stellar absorption feature. We use CO 2-0 2.29µm in the K-band or
CO 6-3 1.62µm in the H-band. Although the equivalent widths WCO2−0 and WCO6−3 vary considerably for
individual stars, the integrated values for stellar clusters span only a rather limited range. This was shown
by Oliva et al. (1995) who measured these values for elliptical, spiral, and star-forming (Hii) galaxies. We
have plotted their measurements of these two absorption features in the left-hand panel of Fig. 2, together
with the equivalent widths of giant and supergiant stars from Origlia et al. (1993).
In STARS we use empirically determined equivalent widths from library spectra (Förster Schreiber
2000) to compute the time-dependent equivalent width for an entire cluster of stars. Results for various star
formation histories are shown in the centre and right panels of Fig. 2, for WCO6−3 and WCO2−0 respectively.
Typical values are WCO6−3 = 4.5Å and WCO2−0 = 12Å. The dashed box in the left panel shows that the
– 7 –
locus of 20% deviation from each of these computed values is consistent with observations. That the Hii
galaxies have slightly higher WCO2−0 can be understood because these are selected to have bright emission
lines and hence are strongly biassed towards young stellar ages – often corresponding to the maximum depth
of the stellar features that occurs at 10 Myr due to the late-type supergiant population. It may be this
bias for galaxies selected as ‘starbursts’, and the similarity of the CO depth for starbursts of all other ages,
that led Ivanov et al. (2000) to conclude that there is no evidence for strong starbursts in Seyfert 2 galaxies.
Similarly, an estimate of the dilution can be found from the Nai 2.206µm line. Fig. 7 of Davies et al. (2005)
shows that for nearly all star formation histories the value WNa I remains in the range 2–3Å.
Our conclusion here is that within a reasonable uncertainty of ±20% (see Fig. 2), one can assume that
the intrinsic equivalent width of the absorption – most notably CO – features of any stellar population that
contains late-type stars is independent of the star formation history and age. For a stellar continuum diluted
by additional non-stellar emission, the fraction of stellar light is
fstellar = Wobs/Wint
where Wobs and Wint are the observed and intrinsic equivalent widths of the CO features discussed above.
Thus, we are able to correct the observed continuum magnitude for the contribution associated with the
3.2. Stellar colour and luminosity
Our data cover both the H and K-bands – hence the reason for using both WCO6−3 and WCO2−0. In
order to homogenize the dataset, we need to convert H-band stellar magnitudes to K-band. The STARS
computation in Fig. 3 shows that this conversion is also independent of the star formation history, being
close to H −K = 0.15 mag (no extinction) for all timescales and ages. This result is supported empirically
by photometry of elliptical and spiral galaxies performed by Glass (1984). For ellipticals H −K ∼ 0.2–0.25,
and for spirals H − K ∼ 0.2–0.3. Some of the difference between the data and models could be due to
extinction since H −K = (H −K)0 + (AH −AK); and for AV = 1, AH −AK = 0.08. However, at the level
of precision required here, the 5–10% difference between model and data can be considered negligible.
To convert from absolute magnitude to luminosity we use the relation
MK = −0.33 − 2.5 logLK
where LK is the total luminosity in the 1.9–2.5µm band in units of bolometric solar luminosity (1L⊙ =
3.8 × 1026 W), and as such different from the other frequently used monochromatic definition with units of
the solar K-band luminosity density (2.15 × 1025 Wµm−1). We then use STARS to estimate the bolometric
stellar luminosity Lbol. The relation between Lbol and LK is shown in the right panel of Fig. 3. The
dimensionless ratio Lbol/LK depends on the age and the exponential decay timescale of the star formation.
However, the range spanned is only 20–200 for ages greater than 10 Myr. Thus even if the star formation
history cannot be constrained, a conversion ratio of Lbol/LK ∼ 60 will have an associated uncertainty of
only 0.3 dex. In general we are able to apply constraints on the star formation age, and so our errors will be
accordingly smaller.
– 8 –
3.3. Specific Star Formation Diagnostics
Graphs showing how the diagnostics vary with age and star formation timescale are shown in Fig. 4.
3.3.1. Brγ equivalent width
Once the stellar continuum luminosity is known, an upper limit to the equivalent width of Brγ asssociated
with star formation can be found from the narrow Brγ line flux. In some cases it is possible to estimate what
fraction of the narrow Brγ might be associated with the AGN. This can be done both morphologically, for
example if the line emission is extended along the galaxy’s minor axis; and/or kinematically, for example if
the line shows regions that are broader, perhaps with FWHM a few hundred km s−1, suggestive of outflow.
Even if acounting for the AGN contribution is not possible, one may be able to set interesting upper limits or
even rule out continuous star formation scenarios, and put a constraint on the time since the star formation
was active. This can be seen in the lefthand panel of Fig. 4, which shows for example that for ages less than
109 yrs, continuous star formation scenarios will always have WBrγ > 12 Å.
3.3.2. Supernova rate
We estimate the type ii (core collapse) supernova rate νSN from the radio continuum using the relation
(Condon 1992):
LN (W Hz
−1) = 1.3 × 1023 ν−α(GHz) νSN (yr
where LN is the non-thermal radio continuum luminosity, ν is frequency of the observation and α ∼ 0.8 the
spectral index of the non-thermal continuum. This relation was derived for Galactic supernova remnants;
but a similar one, differing only in having a coefficient of 1.1 × 1023, was derived by Huang et al. (1994) for
M 82. For the 5 GHz non-thermal radio continuum luminosity of Arp 220 (176 mJy, Anantharamaiah et al.
2000) it would lead to a supernova rate of 2.9 yr−1, comfortably within the 1.75–3.5yr−1 range estimated by
Smith et al. (1998) based on the detection of individual luminous radio supernovae. This, therefore, seems
a reasonable relation to apply to starbursts.
We have to be careful, however, to take into account any contribution from the AGN to the radio
continuum. Our premise for the nuclei of Seyfert galaxies is that if the nuclear radio continuum is spatially
resolved (i.e. it has a low brightness temperature) and does not have the morphology of a jet, it is likely
to originate in extended star formation. At the spatial scales of a few parsec or more that we can resolve,
emission from the AGN will be very compact. As a result, we can use the peak surface brightness to
estimate the maximum (unresolved) contribution from an AGN. Wherever possible, we use radio continuum
observations at a comparable resolution to our data to derive the extended emission; and observations at
higher resolution to estimate the AGN contribution. Details of the data used in each case are given in
the relevant sub-sections for each object in Appendix A. In addition, we exclude any emission obviously
associated with jets, for example as in NGC 1068.
To use νSN as a diagnostic, we normalise it with respect to the stellar K-band luminosity. This gives
the ratio 1010νSN/LK , for which STARS output is drawn in Fig. 4.
– 9 –
3.3.3. Mass-to-light ratio
Models indicate that the ratio M/LK of the stellar mass to K-band luminosity should be an excellent
diagnostic since, for ages greater than 10 Myr, it increases monotonically with age as shown in Fig. 4.
However, in practice estimating the stellar mass is not entirely straightforward. In many cases it is only
practicable to derive the dynamical mass. It may be possible to estimate and hence correct for the molecular
gas mass based on millimetre CO maps, but these are scarce at sufficiently high spatial resolution and are
associated with their own CO-to-H2 conversion uncertainties. We also note that it is often not possible to
separate the ‘old’ and ‘young’ stellar populations. The best one can do is estimate the overall mass-to-light
ratio, and argue that this is an upper limit to the true ratio for the young population. While there inevitably
remains uncertainty on the true ratio, the limit is often sufficient to apply useful constraints on the age of
the ‘young’ population.
Our estimates of the dynamical mass are based wherever possible on the stellar kinematics, since the gas
kinematics can be perturbed by warps, shocks, and outflows. We begin by estimating the simple Keplerian
mass assuming that the stars are supported by ordered rotation at velocity Vrot = Vobs/ sin i in a thin plane.
However, the stellar kinematics in all the galaxies exhibit a significant velocity dispersion indicating that a
considerable mass is supported by random rather than ordered motions. Thus the simple Keplerian mass
is very much an underestimate, and any estimate of the actual mass is associated with large uncertainties
– see for example Bender et al. (1992), who derive masses of spheroidal systems. As stated in Section 2,
we assume that the random motions are isotropic. Our relation for estimating the mass enclosed within a
radius R is then
M = (V 2rot + 3σ
2)R/G.
where σ is the observed 1-dimensional velocity dispersion.
We note that when taking rotation into account in estimating the masses of spheroids with various
density profiles, Bender et al. (1992) also use a factor 3 between the V and σ terms in their Appendix B.
Despite the complexities involved, within the unavoidable uncertainties (a factor 2–3), their relation gives
the same mass as that above. Although this uncertainty appears to be quite large, it does not impact the
results and conclusions in this paper since we are concerned primarily with order-of-magnitude estimates
when considering mass surface densities.
4. Properties of Nuclear Star Formation
In the following section we bring to together the individual results (detailed in Appendix A) to form a
global picture. It is possible to do this because all the data have been analysed in a consistent manner, using
the tools described in Section 3 to compare in each object the same diagnostics to the same set of stellar
evolutionary synthesis models.
We note that the discussion that follows is based on results for 8 of the AGN we have observed. As
explained in Appendix A we exclude NGC 2992 because we are not able to put reliable constraints on the
properties of the nuclear star formation. Despite this, there are indications that at higher spatial resolution
one should expect to find a distinct nuclear stellar population as has been seen in other AGN.
– 10 –
Size scale Tracing the stellar features rather than the broad-band continuum, we have in all cases resolved
a stellar population in the nucleus close around the AGN. While this should not be unexpected if the stellar
distribution follows a smooth r1/4 or exponential profile, we have in several cases been able to show that on
scales of < 50 pc there is in fact an excess above what one would expect from these profiles. This suggests
that in general we are probing an inner star forming component.
Fig. 5 shows normalised azimuthally averaged stellar luminosity profiles for the AGN. These have not
been corrected for a possible old underlying population, nor has any deconvolution with the PSF been
performed. Nevertheless, it is still clear that the stellar intensity increases very steeply towards the nucleus.
In 6 of the 8 galaxies shown, the half-width at half-maximum is less than 50 pc. The remaining 2 galaxies
are the most distant in the sample, and the spatial resolution achieved does not permit a size measurement
on these scales. We may conclude that the physical radial size scale of the nuclear star forming regions in
Seyfert galaxies does not typically exceed 50 pc.
Stellar Age For 8 of the AGN studied here, we have been able to use classical star formation diagnostics
based on line and continuum fluxes as well as kinematics to constrain the ages of the inner star forming
regions. The resulting ages should be considered ‘characteristic’, since in many cases there may simultane-
ously be two or more stellar populations that are not co-eval. For example, if a bulge population exists on
these small spatial scales, it was not usually possible to account for the contamination it would introduce.
While this would have little effect on WBrγ , it could impact on M/LK more strongly, increasing the inferred
age. The ages we find lie in the range 10–300Myr, compelling evidence that it is common for there to be
relatively young star clusters close around AGN.
Intriguingly, we also find rather low values of WBrγ : typically WBrγ . 10 Å (see Table 3). This indicates
directly that there is currently little or no on-going star formation. Coupled with the relatively young ages,
we conclude that the star formation episodes are short-lived. One may speculate then that the star formation
is episodic, recurring in short bursts. The scale of the bursts and time interval between them would certainly
have an impact on the fraction of Seyfert nuclei in which observational programmes are able to find evidence
for recent star formation.
Nuclear Stellar Disks The first evidence for nuclear stellar disks came from seeing limited optical
spectroscopy, for which a slight reduction in σ∗ was seen for some spiral galaxies (Emsellem et al. 2001;
Márquez et al. 2003; Shapiro et al. 2003). And there is now a growing number of spiral galaxies – more than
30 – in which the phenomenon has been observed, suggesting that they might occur in 30% or more of disk
galaxies (Emsellem 2006c). The σ∗-drop has been interpreted by Emsellem et al. (2001) as arising from a
young stellar population that is born from a dynamically cold gas component, and which makes a significant
contribution to the total luminosity. This appears to be borne out by N-body and SPH simulations of iso-
lated galaxies (Wozniak et al. 2003), which suggest that although the entire central system will slowly heat
up with time, the σ∗-drop can last for at least several hundred Myr. Indeed, preliminary analysis of optical
integral field data for NGC 3623 suggest that the stellar population responsible for the σ∗-drop cannot be
younger than 1 Gyr (Emsellem 2006b).
Our results provide strong support for the nuclear disk interpretation. In previous work (Davies et al.
2006; Mueller Sánchez et al. 2006), we had argued that in both Circinus and NGC 3227 the inner distributions
were disk-like, albeit thickened. We have now found much more direct evidence for this phenomenon in
NGC 1097 and NGC 1068. In both of these galaxies, we have spatially resolved a σ∗-drop and an excess
– 11 –
stellar continuum over the same size scales. In NGC 1097 this size was ∼0.5′′, corresponding to about 40 pc.
For NGC 1068 these effects were measured out to ∼1′′, equivalent to 70 pc. These are not the scale lengths
of the disks, but simply the maximum radius to which we can detect them. In both cases the mean mass
surface densities are of order Σ =(1-3)×104 M⊙ pc
−2. For an infinitely large thin self gravitating stellar disk,
one can use the expression σ2z = 2πGΣz0 to estimate the scale height. Although this may not be entirely
appropriate, we use it here to obtain a rough approximation to the scale heights, which are 5–20 pc. Thus
while the disks appear to be flattened, they should still be considered thick since the radial extent is only a
few times the scale height.
The impact of nuclear starbursts on the central light profile of galaxies was considered theoretically more
than a decade ago by Mihos & Hernquist (1994). They performed numerical simulations of galaxy mergers
to study the mass and luminosity profiles of the remnants, taking gas into account, and estimating the star
formation rate using a modified Schmidt law. They found that there should be a starburst in the nucleus
which would give rise to an excess stellar continuum above the r1/4 profile of the older stars in the merged
system. Several years ago, compact nuclei were found to be present in a significant fraction of spiral galaxies
(Balcells et al. 2003) as well as Coma cluster dwarf ellipticals (Graham & Guzmán 2003). More recently,
nuclei with a median half-light radius of 4.2 pc have been found in the majority of early-type members of
the Virgo Cluster (Côté et al. 2006); and traced out to ∼1′′, equivalent to ∼ 100 pc, in some of the ‘wet’
merger remnants in that cluster (Kormendy et al. 2007). While the nuclear starbursts in these latter cases
are caused by a merger event, whereas those we are studying arise from secular evolution as gas from the
galaxy disk accretes in the nucleus, there appear to be many parallels in the phenomenology of the resulting
starbursts.
Star Formation Rate It is possible to estimate the bolometric luminosity Lbol∗ of the stars from their K-
band luminosity LK even if one knows nothing about the star formation history. As discussed in Section 3,
this would result in an uncertainty of about a factor 3. The diagnostics in Table 3 and discussions in
Appendix A enable us to apply some constraints to the characteristic age of the star formation. Because
continuous star formation is ruled out by the low WBrγ , we have assumed exponential decay timescales of
τSF = 10–100Myr. We have then used STARS to estimate the average star formation rates. In order to
allow a meaningful comparison between the objects, the rates have been normalised to the same area of
1 kpc2. These are the rates given in Table 3. They are calculated simply as the mass of stars produced
divided by the entire time since the star forming episode began. Because τSF is shorter than the age, the
average includes both active and non-active phases of the starburst. Indeed, for τSF = 10 Myr one would
expect the star formation rate during the active phases to be at least a factor of a few, and perhaps an
order of magnitude, greater. The table shows that on scales of a few hundred parsecs one might expect a
few ×10 M⊙ yr
−1 kpc−2, while on scales of a few tens of parsecs mean rates reaching ∼ 100 M⊙ yr
−1 kpc−2
should not be unexpected; and correspondingly higher – up to an order of magnitude, see Fig. 6 – during
active phases.
An obvious question is why there should be such vigorous star formation in these regions. Star formation
rates of 10–100M⊙ yr
−1 kpc−2 are orders of magnitude above those in normal galaxies and comparable to
starburst galaxies. The answer may lie in the Schmidt law and the mass surface densities we have estimated
in Table 3. Fig. 7 shows these surface densities at the radii over which they were estimated, revealing a trend
towards higher densities on smaller scales and values of a few times 104 M⊙ pc
−2 in the central few tens of
parsecs. The global Schmidt law, as formulated by Kennicutt (1998), states that the star formation rate
depends on the gas surface density as ΣSFR ∝ Σ
gas. If one assumes that 10–30% of the mass in our AGN is
– 12 –
gas, then this relation would predict time-averaged star formation rates in the range 10–100M⊙ yr
−1 kpc−2,
as have been observed. That the high star formation rates may simply be a consequence of the high mass
surface densities is explored futher by Hicks et al. (in prep.).
Stellar Luminosity As a consequence of the high star formation rates, the stellar luminosity per unit area
close around the AGN is very high in these objects. Despite this, because the star formation is occurring
only in very small regions, the absolute luminosities are rather modest. This can be seen in Fig. 8 which
shows the bolometric luminosity of the stars as a fraction of the entire bolometric luminosity of the galaxy.
We have calculated a range for the ratio Lbol∗/LK appropriate for each galaxy based on the ages in Table 3
for different τSF. Because we assume that all the K-band stellar continuum is associated with the young
stars, we have adopted the lower end of each range in an attempt to minimise possible overestimation of
Lbol∗. The resulting values for the ratio used span 30–130, within a factor of 2 of the ‘baseline’ value of 60
given in Section 3. In the central few tens of parsecs, young stars contribute a few percent of the total. But
integrated over size scales of a few hundred parsecs, this fraction can increase to more than 20%. On these
scales, the star formation is energetically significant when compared to the AGN. Such high fractions imply
that on the larger scales the extinction to the young stars must be relatively low. On the other hand, on the
smallest scales where in absolute terms the stellar luminosity is small, there could in general be considerable
extinction even at near infrared wavelengths. In this paper we have not tried to account for extinction since
it is very uncertain. The primary effect of doing so would simply be to increase the stellar luminosity above
the values discussed here.
Fig. 9 shows the stellar bolometric luminosity Lbol∗ integrated as a function of radius. All the curves
follow approximately the same trend, with the luminosity per unit area increasing towards smaller scales and
approaching 1013 L⊙ kpc
−2 in the central few parsecs. This appears to be a robust trend and will not change
significantly even with large uncertainties of a factor of a few. It is remarkable that the luminosity density
of 1013 L⊙ kpc
−2 is that estimated by Thompson et al. (2005) for ULIRGs, which they modelled as optically
thick starburst disks. The main difference between the ULIRG model and the starbursts close around AGN
is the spatial scales on which the starburst occurs.
Based on this model, they argued that ULIRGs are radiating at the Eddington limit for a starburst,
defined as when the radiation pressure on the gas and dust begins to dominate over self-gravity. The limiting
luminosity-to-mass ratio was estimated to be ∼ 500 L⊙/M⊙ by Scoville (2003). He argued that in a star
cluster, once the upper end of the main sequence was populated, the radiation pressure would halt further
accretion on to the star cluster and hence terminate the star formation. Following Thompson et al. (2005),
we apply this definition to the entire disk rather than a single star cluster. For the 1013 L⊙ kpc
−2, this implies
a mass surface density of 2 × 104 M⊙ pc
−2. Comparing these quantities to the AGN we have observed, we
find that on scales of a few tens of parsecs they are an order of magnitude below the Eddington limit. On
the other hand, we have already seen that the low WBrγ indicates that there is little on-going star formation
and hence that the starbursts are short-lived. This is important because short-lived starbursts fade very
quickly. As shown in Fig. 6, for a decay timescale of τSF = 10 Myr, Lbol∗ will have decreased from its peak
value by more than an order of magnitude at an age of 100 Myr. Thus it is plausible – and probably likely
– that while the star formation was active, the stellar luminosity was an order of magnitude higher. In this
case the starbursts would have been at, or close to, their Eddington limit at that time.
The luminosity-to-mass ratio of 500 L⊙/M⊙ associated with the Eddington limit is in fact one that all
young starbursts would exceed if, beginning with nothing, gas was accreted at the same rate that it was
converted into stars. That, however, is not a realistic situation. A more likely scenario, shown in Fig. 10, is
– 13 –
that the gas is already there in the disk. In this case, a starburst with a star-forming timescale of 100 Myr
could never exceed 100 L⊙/M⊙. To reach 500 L⊙/M⊙, the gas would need to be converted into stars on a
timescale . 10 Myr. This timescale is independent of how much gas there is. Thus, for a starburst to reach
its Eddington limit, it must be very efficient, converting a significant fraction of its gas into stars on very
short ∼ 10 Myr timescales. This result is consistent with the prediction of the Schmidt law, which states that
disks with a higher gas surface density will form stars more efficiently. The reason is that the star formation
efficiency is simply SFE = ΣSFR/Σgas ∝ Σ
gas. Thus, from arguments based solely on the Schmidt law and
mass surface density, one reaches the same conclusion that the gas supply would be used rather quickly and
the lifetime of the starburst would be relatively short.
Summarising the results above, a plausible scenario could be as follows. The high gas density leads to
a high star formation rate, producing a starburst that reaches its Eddington limit for a short time. Because
the efficiency is high, the starburst can only be active for a short time and then begins to fade. Inevitably,
one would expect that the starburst is then dormant until the gas supply is replenished by inflow. This
picture appears to be borne out by the observations presented here.
5. Starburst-AGN Connection
In the previous sections we have presented and discussed evidence that in general there appears to have
been moderately recent star formation on small spatial scales around all the AGN we have observed. Fig. 11
shows the first empirical indication of a deeper relationship between the star formation and the AGN. In this
figure we show the luminosity of the AGN, both in absolute units of solar luminosity and also in relative units
of its Eddington luminosity LEdd, against the age of the most recent known nuclear star forming episode.
Since the AGN luminosity is not well known, we have made the conservative assumption that it is equal to
half the bolometric luminosity of the galaxy – as may be the case for NGC 1068 (Pier et al. 1994, but see
also Bland-Hawthorne et al. 1997). To indicate the expected degree of uncertainty in this assertion we have
imposed errorbars of a factor 2 in either direction, equivalent to stating that the AGN luminosity in these
specific objects is likely to be in the range 25–100% of the total luminosity of the galaxy. The Eddington
luminosity is calculated directly from the black hole mass, for which estimates exist for these galaxies from
reverberation mapping, the MBH − σ∗ relation, maser kinematics, etc. These are listed in Table 2. For the
age of the star formation, we have plotted the time since the most recent known episode of star formation
began, as given in Table 3. For galaxies where a range of ages is given, we have adopted these to indicate
the uncertainty; the mean of these, ∼ ±30%, has been used to estimate the uncertainty in the age for the
rest of the galaxies. We note that these errorbars reflect uncertainties in characterising the age of the star
formation from the available diagnostics and also in the star formation timescale τSF. However, there are still
many implicit assumptions in this process, and we therefore caution that the actual errors in our estimation
of the starburst ages may be larger than that shown.
Conceding this, we do not wish to over-interpret the figure. Keeping the uncertainties in mind, Fig. 11
shows the remarkable result that AGN which are radiating at lower efficiency . 0.1 L/LEdd are associated
with younger . 50–100Myr starbursts; while those which are more efficient & 0.1 L/LEdd have older &
50–100Myr starbursts. If one were to add to this figure the Galactic Centre – which is known to have
an extremely low luminosity (L/LEdd < 10
−5; Ozernoy & Genzel 1996; Baganoff et al. 2003) and to have
experienced a starburst 6 ± 2 Myr ago (Paumard et al. 2006) – it would be consistent with the categories
above. The inference is that either there is a delay between the onset of starburst activity and the onset of
AGN activity, or star formation is quenched once the black hole has become active.
– 14 –
In Section 4 we argued that the starbursts are to some extent self-quenching: that very high star forma-
tion efficiencies are not sustainable over long periods. In addition, an intense starburst will provide significant
heat input to the gas, which is perhaps partially responsible for the typically high gas velocity dispersions
in these regions (Hicks et al., in prep.). This itself could help suppress further star formation. Heating by
the AGN could also contribute to this process, and has been proposed as the reason why the molecular
torus is geometrically thick (Pier & Krolik 1992; Krolik 2007). It is also used to modulate star formation (at
least on global scales) in semi-analytic models of galaxy evolution (Granato et al. 2004; Springel et al. 2005).
While this is certainly plausible, it does not explain either why the star formation in some galaxies with a
lower luminosity AGN has already ceased, nor why none of the AGN associated with younger starbursts are
accreting efficiently.
Instead we argue for the former case above, that efficient fuelling of a black hole is associated with a
starburst that is at least 50–100Myr old. It may be because of such a delay between AGN and starburst
activity that recent star formation is often hard to detect close to AGN: the starburst has passed its most
luminous (very young) age, and is in decline while the AGN is in its most active phase (see Fig. 6). This
does not necessarily imply that the a priori presence of a starburst is required before an AGN can accrete
gas – although it seems inevitable that one will occur as gas accumulates in the nucleus. Nor does it imply
that all starbursts will result in fuelling a black hole; indeed it is clear that there are many starbursts not
associated with AGN. As we argue below, the crucial aspect may be the stellar ejecta associated with the
starburst; and in particular, not just the mass loss rate, but the speed with which the mass is ejected.
Winds from OB stars In the Galactic Center, Ozernoy & Genzel (1996) proposed that it is the recent
starburst there that is limiting the luminosity of the black hole. In this scenario, mechanical winds from
young stars – both the outflow and the angular momentum of the gas (which is a consequence of the angular
momentum of the stars themselves) – hinder further inflow. The authors argued that almost none of the
gas flowing into the central parsec reached the black hole because of outflowing winds from IRS 16 and He i
stars in that region. Detailed modelling of the Galactic Center region as a 2-phase medium was recently
performed by Cuadra et al. (2006). They included both the fast young stellar winds with velocities of
700 km s−1 (Ozernoy et al. 1997) and the slower winds of ∼ 200 km s−1 (Paumard et al. 2001); and also took
into account the orbital angular momentum of the stars (Paumard et al. 2001; Genzel et al. 2003), which
had a strong influence on reducing the accretion rate. They found that the average accretion rate onto
the black hole was only ∼ 3 × 10−6 M⊙ yr
−1, although an intermittent cold flow superimposed considerable
variability onto this. In contrast, the hypothetical luminosity Ozernoy & Genzel (1996) estimate that Sgr A∗
would have if it could accrete all the inflowing gas, would be 5 × 1043 erg s−1, typical of Seyfert galaxies.
In principle this process could be operating in other galaxy nuclei where there has been a starburst which
extends to less than 1 pc from the central black hole. However, it cannot explain the timescale of the delay we
have observed, which is an order of magnitude greater than the main sequence lifetime of OB and Wolf-Rayet
stars.
Winds from AGB stars Stars of a few (1–8 M⊙) solar masses will evolve on to the asymptotic giant
branch (AGB) at the end of their main sequence lifetimes. The timescale for stars at the upper end of this
range to reach this phase is ∼ 50 Myr, comparable to the delay apparent in Fig. 9. Since AGB stars are known
to have high mass-loss rates, of order 10−7–10−4 M⊙ yr
−1 at velocities of 10–30km s−1 (Winters et al. 2003),
they may be prime candidates for explaining the delay between starburst and AGN activity. To quantify
this, we consider how much of the mass in the wind could be accreted by the central supermassive black
– 15 –
hole.
The Bondi parameterisation of the accretion rate onto a point particle for a uniform spherically sym-
metric geometry is given by (Bondi 1952)
2πG2 M2 ρ
(V 2 + c2s)
where M is the mass of the point particle moving through a gas cloud, V is the velocity of the particle with
respect to the cloud, ρ is the density of the cloud far from the point particle, and cs is the sound speed.
This approximation is still used to quantify accretion on to supermassive black holes in models of galaxy
evolution (Springel et al. 2005), even though it may be significantly inaccurate for realistic (e.g. turbulent)
media (Krumholz et al. 2006). Here, it is sufficient to provide an indication of the role that stellar winds
may play in accretion onto a central black hole. The density of the stellar wind at a distance R from the
parent star is given by
ρwind =
Ṁwind
4R2 Vwind
In our case, R is the distance from the star to the black hole. One would therefore expect that the accretion
rate on to the black hole could be written as (see also Melia 1992)
ṀBH ∼
G2 M2BH Ṁwind
+ c2s)
3/2 Vwind R2
This equation shows that ṀBH ∝ V
. We have implicitly assumed that Vwind is greater than the orbital
velocity Vorb of the star from which it originates. This is not the case for AGB winds, and so one reaches
the limiting case of ṀBH ∝ V
, where for the galaxies we have observed Vorb ∼ 50–100km s
−1. This is still
at least an order of magnitude less than the winds from OB and Wolf-Rayet stars. Thus, even though the
mass loss rates from individual OB and Wolf-Rayet stars are similar to those of AGB stars, the AGB winds
will fuel a black hole much more efficiently. However, for slow stellar winds that originate close to a 107 M⊙
black hole, the equation breaks down because the conditions of uniformity and spherically symmetry are
strongly violated. Indeed, the apparent accretion rate exceeds the outflow rate – implying that essentially
the entire wind can be accreted. For AGB wind velocities of 10–30 km s−1, the maximum radius at which
the entire wind from a star in Keplerian orbit around a 107 M⊙ black hole will not exceed the escape velocity
from that orbit (i.e. Vwind + Vorb < Vesc) is around 10–70pc. We adopt the middle of this range, 40 pc, as
the characteristic radius within which it is likely that a significant fraction, and perhaps most, of the AGB
winds are accreted onto the black hole. Fig. 9 indicates that the stellar luminosity within this radius is
∼ 2 × 109 L⊙. It is this luminosity that has been used to scale the STARS model (for τSF = 10 Myr and an
age of 100 Myr) in Fig. 6, and so one can also simply read off the mass loss from the figure. The mass loss
rate for such winds peaks at about 0.1 M⊙ yr
−1 and then tails off proportionally to the K-band luminosity,
leading to a cumulative mass lost of 2× 107 M⊙ after 1 Gyr (although most of the loss occurs actually occurs
within half of this timespan). This mass loss rate is sufficient to power a Seyfert nucleus for a short time. A
typical Seyfert with MBH ∼ 10
7 M⊙ requires 0.02 M⊙ yr
−1 to radiate at the Eddington limit. Even for the
short bursts we have modelled, Fig. 6 shows that this can be supplied by AGB winds for starburst ages in
the range 50–200Myrs.
We note that taking an AGB star luminosity of 104 L⊙ (which is at the high end of the likely average,
Nikolaev & Weinberg 1997) we then find that there are ∼ 2 × 105 AGB stars close enough to the black hole
to contribute to accretion. In order to provide at least 0.02 M⊙ yr
−1, the typical mass loss rate per star must
– 16 –
exceed 10−7 M⊙ yr
−1, which is the lower limit of the range measured for Galactic AGB stars given above.
Thus the mass losses and rates estimated here appear to be plausible.
The low speed of these winds means they will not create much turbulence. We quantify this by consid-
ering their total mechanical energy 1
mv2 integrated over the same timespan, which is ∼ 1045 J. These two
quantities – gas mass ejected and mechanical energy – are compared to those for supernovae below.
Supernovae Type ii supernovae are the stellar outflows most able to create turbulence in the interstellar
medium, since they typically eject masses of ∼ 5 M⊙ at velocities of ∼ 5000 km s
−1 (Chevalier 1977). Each
supernova therefore represents a considerable injection of mechanical momentum and energy into the local
environment. A large number of compact supernova remnants are known, for example in M 82 and Arp 220,
and are believed to have expanded into dense regions with nH ∼ 10
3–104 cm−3 (Chevalier & Fransson 2001).
These authors argue that such remnants become radiative when they reach sizes of ∼ 1 pc, at which point
the predicted expansion velocity will have slowed to ∼ 500 km s−1. By this time, the shock front will have
driven across ∼ 1000 M⊙ of gas. When integrated over the age of the starburst, even for low supernova
rates – e.g. the current rate within 30 pc of the nucleus of NGC 3227 is ∼ 0.01 yr−1 (Davies et al. 2006)
– this represents a substantial mass of gas that has been affected by supernova remnants. The STARS
model we have constructed in Fig. 6 indicates that typically one could expect ∼ 106 supernovae to occur as
a result of one of the short-lived starbursts; and that most of these will occur around 10–50Myr after the
beginning of the starburst. For a decay timescale of the star formation rate that is longer than τSF = 10 Myr,
this timespan will increase. Hence, supernovae may also play a role in causing the observed delay between
starburst and AGN activity.
STARS calculates the mass loss and mass loss rates using a very simple scheme, assuming that a star
ejects all of its lost mass at the end of its life on a stellar track. Thus, it does not calculate the mass lost
from supernovae explicitly, rather the combined mass lost from OB winds and supernovae which is much
higher. We therefore adopt the ∼ 5 M⊙ per supernova given above, which yields a total ejected mass of
∼ 8 × 106 M⊙. This is about 40% of that released by AGB winds. However, since this gas is ejected at high
speed and ṀBH ∝ V
, the efficiency with which it can be accreted onto the black hole is extremely low.
This can also be seen in the total mechanical energy of ∼ 1050 J, which is several orders of magnitude greater
than for AGB winds. In fact the total mechanical energy exceeds the binding energy of the nuclear region,
which is of order 1048 J (assuming 108 M⊙ within 40 pc). As a result, it is highly likely that supernova cause
some fraction of the gas to be permanently expelled. Indeed, superwinds driven by starbursts are well known
in many galaxies. This is not important as long as sufficient gas either remains to fuel the AGN, or more is
produced by stellar winds – which, as we have argued above, appears to be the case for AGB stars.
6. Conclusions
We have obtained near infrared spectra of 9 nearby active galactic nuclei using adaptive optics to achive
high spatial resolution (in several cases better than 10 pc). For 7 of these, integral field spectroscopy with
SINFONI allows us to reconstruct the full 2-dimensional distributions and kinematics of the stars and gas.
Although the individual AGN are very varied, we have analysed them in a consistent fashion to derive: the
stellar K-band luminosity, the dynamical mass, and the equivalent width of the Brγ line. We have combined
these with radio continuum data from the literature, which has been used to estimate the supernova rate.
We have used these diagnostics to constrain STARS evolutionary synthesis models and hence characterize
– 17 –
the star formation timescales and ages of the starbursts close around AGN. Our main conclusions can be
summarised as follows:
• The stellar light profiles show a bright nuclear component with a half-width at half-maximum of
less than 50 pc. In a number of cases these nuclear components clearly stand out above an inward
extrapolation of the profile measured on larger scales. In addition, there are 2 cases which show
kinematical evidence for a distinct stellar component, indicating that the nuclear stellar populations
most probably exist in thick nuclear disks. The mean mass surface densities of these disks exceeds
104 M⊙ pc
• There is abundant evidence for recent star formation in the last 10–300Myr. But the starbursts are no
longer active, implying that the star formation timescale is short, of order a few tens of Myr. While the
starbursts were active, the star formation rates would have been much higher than the current rates,
reaching as high as 1000 M⊙ kpc
−2 in the central few tens of parsecs (comparable to ULIRGs, but on
smaller spatial scales). These starbursts would have been Eddington limited. Due to the very high star
forming efficiency, the starbursts would have also exhausted their fuel supply on a short timescale and
hence have been short-lived. It therefore seems likely that nuclear starbursts are episodic in nature.
• There appears to be a delay of 50–100Myr (and in some cases perhaps more) between the onset of star
formation and the onset of AGN activity. We have interpreted this as indicating that the starburst
has a significant impact on fuelling the central black hole, and have considered whether outflows from
stars might be responsible. While supernovae and winds from OB stars eject a large mass of gas, the
high velocity of this gas means that its accretion efficiency is extremely low. On the other hand, winds
from AGB stars ultimately dominate the total mass ejected in a starburst; and the very slow velocities
of these winds mean they can be accreted onto the black hole very efficiently.
The authors thank all those who assisted in the observations, and also the referee for a thorough review
of the paper. This work was started at the Kavli Institute for Theoretical Physics at Santa Barabara and
as a result was supported in part by the National Science Foundation under Grant No. PHY05-51164. RD
aknowledges the interesting and useful discussions he had there with Eliot Quartaert, Norm Murray, Julian
Krolik and Todd Thompson.
Facilities: Keck:II (NIRSPAO, NIRC2), VLT:Yepun (NACO, SINFONI).
A. Individual Objects
This appendix contains specific details on the individual objects. We summarize our published results
from near infrared adaptive optics spectroscopy of individual objects, and present a brief analysis of the
new data for several other objects. The aim of re-assessing the data for Mkn 231 that has already been
published is to ensure that it is analysed using STARS in a manner that is consistent with the new data. For
NGC 7469, we make a significant update of the analysis using new data from integral field spectroscopy. In
general, for objects with new data, we provide only the part of the analysis relevant to understanding star
formation around the AGN. Our intention is that a complete analysis for each object will be presented in
future publications.
– 18 –
Our analyses are restricted to the nuclear region. Since there is no strict universal definition of what
comprises the ‘nuclear region’, we explicitly state in Table 3 the size of the region we study in each galaxy.
The table also presents a summary of the primary diagnostics. The way in which these have been derived,
and their likely uncertainties, has been discussed in some detail already in Section 3. As such, the description
of these methods is not repeated, and in this Section we discuss only issues that require special attention.
A.1. Summary of Star Forming Properties of Galaxies already Studied
A.1.1. Mkn 231
A detailed analysis of the star formation in the nucleus of Mkn 231 at a resolution of about 0.18′′ (150 pc)
was given in Davies et al. (2004b). Here we summarize only the main points; no new data is presented, but
the analysis is updated using STARS to make it consistent with the other objects studied in this paper.
The presence of stellar absorption features across the nucleus demonstrates the existence of a significant
population of stars. The radial distribution and kinematics indicate they lie, like the gas (Downes & Solomon
1998), in a nearly face-on disk. Davies et al. (2004b) found that the dynamical mass imposed a strong
constraint on the range of acceptable starburst models, yielding an upper limit to the age of the stars of
around 120 Myr. Re-assessing the mass-to-light ratio using STARS models suggests that for the increased
mass required by a more face-on orientation (i = 10◦) an upper age of 250 Myr is also possible, depending
on the star formation timescale. However, either a small change of only a few degrees to the inclination (to
i = 15◦), or a relatively short star formation timescale of 10 Myr would reduce the limit to the ∼100 Myr
previously estimated. This is more consistent with the extremely high supernova rate.
The stellar luminosity, found from the dilution of the CO absorption (Davies et al. 2004b), indicates
that stars within 1′′ (800 pc) of the nucleus contribute 25–40% of the bolometric luminosity of the galaxy.
Similarly, within 200 pc, stars comprise 10–15% of Lbol. The age, star formation rate, and size scale (disk
scale length of 0.18–0.2′′) are all consistent with high resolution radio continuum imaging (Carilli et al. 1998).
A.1.2. Circinus
Star formation in the central 16 pc of Circinus was addressed by Mueller Sánchez et al. (2006). The
diagnostics given in Table 3 are taken from this reference. We used the depth of the CO 2-0 bandhead to
estimate the stellar luminosity, combined with the narrow Brγ flux (which we argued originated in star
forming regions rather than the AGN narrow line region) and the radio continuum, to constrain starburst
models. The conclusion was that the starburst was less than 80 Myr old and was already decaying. On
these scales it contributes 1.4% of Lbol, or more if extinction is considered. A similar nuclear star formation
intensity was estimated by Maiolino et al. (1998), who were also able to study Circinus on larger scales.
They found that the luminosity of young stars within 200 pc of the AGN was of order 1010L⊙, and hence
comparable to the AGN.
– 19 –
A.1.3. NGC 3227
An analysis similar to that for Circinus was performed on NGC 3227 by Davies et al. (2006), and the
diagnostics given in Table 3 are taken from this reference. In this case we were able to make estimates of
and correct for contributions of: (1) the narrow line region to Brγ, because there were clear regions along
the minor axis that had higher dispersion; (2) the AGN to the radio continuum, by estimating the maximum
contribution from an unresolved source; and (3) the bulge stars to the stellar luminosity, by extrapolating
the radial profile of the bulge to the inner regions. The STARS models yielded the result that in the nucleus,
star formation began approximately 40 Myr ago and must have already ceased. At the resolution of 0.085′′,
the most compact component of stellar continuum had a measured FWHM of 0.17′′, suggesting an intrinsic
size scale of ∼ 12 pc. Young stars within 30 pc of the AGN (i.e. more than just the most compact region)
have a luminosity of ∼ 3 × 109L⊙, which is ∼ 20% of the entire galaxy.
A.2. Star Forming Properties of Galaxies with New Data
A.2.1. NGC 7469
Star formation on large scales in NGC 7469 has been studied by Genzel et al. (1995). They found that
within 800 pc of the nucleus, a region that includes the circumnuclear ring, the luminosity from young stars
was ∼ 3 × 1011L⊙, about 70% of the galaxy’s bolometric luminosity. This situation is similar to that in
Mkn 231. On smaller scales, the nuclear star formation in NGC 7469 was directly resolved by Davies et al.
(2004a) on a size scale of 0.15–0.20′′ (50–65 pc) FWHM. An analysis of the longslit data, similar to that
for Mkn 231, was made – making use of stellar absorption features, kinematics, and starburst models. We
estimated that the age of this region was no more than 60 Myr under the assumption that the fraction of
stellar light in the K-band in the central 0.2′′ was 20–30%. Our new integral field SINFONI observations of
NGC 7469 at a spatial resolution of 0.15′′ (measured from both the broad Brγ and the non-stellar continuum
profiles, see Section 2) are used here to make a more accurate estimate of the nuclear K-band luminosity.
They enable us to provide a short update to the detailed analysis in Davies et al. (2004a).
The SINFONI data show that the equivalent width of the 2.3µm CO 2-0 is WCO2−0 = 1.8Å in a 0.8
aperture and 0.9Å in a 0.2′′ aperture. The corresponding K-band magnitudes are K = 10.4 and K = 11.8
respectively. If one takes the intrinsic equivalent width of the 2.3µm CO 2-0 bandhead to be 12Å (see
Section 3), one arrives at a more modest value of 8% for the stellar fraction of K-band continuum in the
0.2′′ aperture. The stellar K-band luminosity in this region is then 6 × 107 L⊙. Comparing this to the
dynamical mass in Davies et al. (2004a) yields a mass-to-light ratio of M/LK ∼ 0.6 M⊙/L⊙. Previously,
extrapolation from a 37 mas slit to a filled aperture had led to an underestimation of the total magnitude
but an overestimation of the stellar contribution. Fortuitously, these uncertainties had compensated each
other. The same analysis for the 0.8′′ aperture yields a K-band stellar luminosity of 3 × 108 L⊙ and hence
M/LK ∼ 1.6 M⊙/L⊙.
The K-band datacube yields estimates of the upper limit to WBrγ of 17Å and 11Å in 0.2
′′ and 0.8′′
apertures respectively. This has been corrected for dilution of the stellar continuum (as described in Section 3)
but not for a possible contribution to the narrow Brγ from the AGN. Hence the actual WBrγ corresponding
to only the stellar line and continuum emission will be less than these values – indicating that the star
formation is unlikely still to be on-going.
We estimate the age of the star formation using the STARS models in Fig. 4. Within the 0.2′′ aperture
– 20 –
this gives 100 Myr, comparable to our original estimate. Such a young age is supported by radio continuum
measurements. With a 0.2′′ beam, Colina et al. (2001) reported that the unresolved core flux in NGC 7469
was 12 mJy at 8.4 GHz. With much higher spatial resolution of 0.03′′, Sadler et al. (1995) reported an upper
limit to the unresolved 8.4 GHz continuum of 7 mJy. We assume that the difference of 5 mJy is due to
emission extended on scales of 10–60 pc which is resolved out of one beam but not the other. As discussed
in Section 3, star formation is a likely candidate for such emission. In this case, we would estimate the
supernova rate to be ∼ 0.1 yr−1 and the ratio 1010 νSN/LK ∼ 3. This is likely to be a lower limit since there
was only an upper limit on the core radio flux density. For a ratio of this order, even allowing for some
uncertainty, Fig 4 implies an age consistent with no more than 100 Myr.
Within the 0.8′′ aperture, which we adopt in Table 3, continuous star formation is inconsistent with
WBrγ . For a star formation timescale of τSF = 100 Myr, the mass-to-light ratio implies an age of 190 Myr,
just consistent with the measured value of WBrγ = 11Å. If some of the narrow Brγ is associated with the
AGN rather than star formation, then a shorter star formation timescale is required. For τSF = 10 Myr, the
ratio M/LK yields an age of 110 Myr.
A.2.2. IRAS 05189-2524
Fig. 12 shows the H-band spectrum integrated across two segments of the NACO slit, located on either
side of the nucleus. It shows that even away from the nucleus, the depth of the stellar absorption features
is only a few percent. We have therefore decomposed the data into the stellar and non-stellar parts using
both the stellar absorption features and the spectral slope of the continuum. The latter method has been
shown to work for well sampled data by Davies et al. (2004a). The rationale is that the hot dust associated
with the AGN will be much redder than the stellar continuum. An AGN component is also expected to be
unresolved for a galaxy at the distance (170 Mpc) of IRAS 05189-2524. The spectral slope was determined
by fitting a linear function to the spectrum at each spatial position along the slit. It is plotted as a function
of position in Fig. 13, showing a single narrow peak. A Gaussian fit to this yields a spatial resolution of 0.12′′
(100 pc) FWHM. The stellar continuum, also shown in Fig. 13, has been determined by summing the four
most prominent absorption features: CO 4-1, Si I, CO 5-2, CO 6-3. While a Gaussian is not an optimal fit to
this profile, it does yield an aproximate size scale, which we find to be 0.27′′ FWHM. Quadrature correction
with the spatial resolution yields an intrinsic size of 0.25′′ (200 pc). As a cross-check, in the figure we have
compared the sum of these two components to the full continuum profile. The good match indicates that
the decomposition appears to be reasonable.
Remarkably, the 200 pc size of the nuclear stellar light is very similar to that of the 8.44 GHz radio con-
tinuum map of Condon et al. (1991). With a beam size of 0.50′′×0.25′′, they resolved the nuclear component
to have an intrinsic size of 0.20′′×0.17′′. In constrast to radio sources which are powered by AGN and have
brightness temperatures Tb ≫ 10
5 K, the emission here is resolved and has a low brightness temperature
of ∼ 4000 K. This implies a star forming origin. Using their scaling relations further suggests that the flux
density corresponds to a supernova rate of ∼ 1 yr−1.
As described in Section 3, we have estimated the stellar luminosity by comparing the H-band spectrum
to a template star to correct for dilution. We used HR 8465 a K 1.5 I star for which the equivalent width
of CO 6-3 is 4.2Å, within the 4–5Å range predicted by STARS in Fig. 2. By extrapolating from the spatial
profiles along the slit we have estimated the integrated equivalent width within a 1.1′′ aperture, for which
Scoville et al. (2000) gave an H-band magnitude of 11.83. Using all four features above we find for the
– 21 –
template W = 14.4Å and for IRAS 05189-2524 WCO6−3 = 6.7Å. This implies that in the central 1.1
approximately 45% of the H-band continuum originates in stars. Using the colour conversion H −K = 0.15
from Fig. 3 (see Section 3) we find a K-band magnitude for the stars of 12.55 mag and hence a K-band stellar
luminosity of 2×109 L⊙. Putting these results together we derive a ratio of supernova rate to K-band stellar
luminosity of νSN[yr
−1]/LK [10
10L⊙] ∼ 5. Applying corrections for extinction and an AGN contribution
would tend to decrease this ratio.
As a second diagnostic we use WBrγ . We estimate the dilution of the K-band continuum via two
methods. Firstly, we measure WNaI = 0.3Å, indicating a stellar fraction of 0.10–0.15. A consistency check is
provided by the H-band dilution, which we extrapolate to the K-band using blackbody functions for the stars
and dust assuming characteristic temperatures of 5000 K and 1000 K respectively. This method suggests the
K-band stellar fraction is around ∼ 0.14. Hence correcting the directly measured equivalent width of the
narrow Brγ for the non-stellar continuum yields WBrγ = 4–5Å.
Since IRAS 05189-2524 is close to face-on (Scoville et al. 2000), it is not straightforward to make a
reliable estimate of the dynamical mass. Nevertheless, requiring νSN/LK to be high while WBrγ is low
already puts significant constraints on the star formation history. Thus, although the star formation has
probably ended, the age is unlikely to be greater than 100 Myr, and could be as low as 50 Myr where νSN/LK
peaks. For such ages the ratio Lbol/LK is in the range 100–150. Hence for the young stars within 0.55
(450 pc) of the nucleus we find Lbol ∼ (2–3) × 10
11 L⊙, about 20% of Lbol for the galaxy.
A.2.3. NGC 2992
The spatial resolution of the K-band data for NGC 2992 has been estimated from both the broad Brγ
and the non-stellar continuum (see Section 2 and 3). The two methods yield symmetric PSFs, with FWHMs
of 0.32′′ and 0.29′′ respectively, corresponding to 50 pc.
The CO 2-0 equivalent width of ∼ 3Å implies a stellar fraction of ∼ 0.25 within a radius of 0.4′′, and
hence a stellar luminosity of LK = 3.5 × 10
7 L⊙.
Unlike IRAS 05189-2524, the radio continuum in NGC 2992 is quite complex. Much of the extended
emission on scales of a few arcsec appears to originate from a superbubble, driven either by the AGN or
by a nuclear starburst. On the other hand, most of the nuclear emission seems to be unresolved. With a
beam size of 0.34′′×0.49′′, Wehrle & Morris (1988) measured the unresolved flux to be 7 mJy at 5 GHz. At
a resolution better than 0.1′′, Sadler et al. (1995) reported a 2.3 GHz flux of 6 mJy. Based on this as well as
non-detections at 1.7 GHz and 8.4 GHz, they estimated the core flux at 5 GHz to be <6 mJy. Taking a flat
spectral index, as indicated by archival data (Chapman et al. 2000), one might expect the 5 GHz core flux
to be not much less than 6 mJy, leaving room for only ∼ 1 mJy in extended emission in the central 0.5′′. If
we assume this difference can be attributed to star formation, it implies a supernova rate of ∼ 0.003 yr−1
and hence 1010 νSN/LK ∼ 1. Fig. 4 shows that a ratio of this order is what one might expect for ages up to
200 Myr. However, given the uncertainty it does not impose a significant constraint.
It is also difficult to quantify what fraction of the narrow Brγ is associated with star formation. This is
made clear in Fig. 14 which shows that the morphology of the line (centre left panel) does not follow that of
the stars (far left). In addition, particularly the south-west side is associated with velocities that are bluer
than the surrounding emission, indicative of motion towards us. The western edge also exhibits high velocity
dispersion. Taken together, these suggest that we may be seeing outflow from the apex of an ionisation cone
– 22 –
with a relatively large opening angle. This interpretation would tend to support the hypothesis that the
radio bubble has been driven by the AGN.
The stellar continuum appears to trace an inclined disk, the north west side of which is more obscured
(Fig. 14). However, the velocity dispersion is high, exceeding 150 km s−1 across the whole field (Fig. 15). This
is similar to the 160 km s−1 reported by Nelson & Whittle (1995) from optical spectroscopy, and suggests
that we are seeing bulge stars. To analyse the radial luminosity profile we have fitted it with both an r1/4
law and exponential profile. The fits in Fig. 16 were optimised at radii r > 0.5′′ and then extrapolated
inwards, convolved with the PSF. Whether one could claim that there is excess continuum in the nucleus
depends on the profile fitted. The r1/4 law provides a stronger constraint since it is more cuspy, and suggests
there is no excess. Although this evidence is inconclusive, Fig. 15 suggests that there is some kinematic
evidence favouring the existence of a distinct nuclear stellar population. This comes in the form of a small
unresolved drop in dispersion at the centre, similar to those in NGC 1097 and NGC 1068. While the evidence
in NGC 2992 is not compelling, the dispersion is consistent with there being an equivalent – but fainter –
nuclear disk on a scale of less than our resolution of 50 pc. In general it seems that the K-band light we are
seeing is dominated by the bulge, and we are therefore unable to probe in detail the inner region where it
seems that more recent star formation has probably occurred.
Thus, although the available data suggest there has likely been recent star formation in the nucleus of
NGC 2992, the only strong constraint we can apply is that continuous star formation in the central arcsec
over the last Gyr can be ruled out since it would require WBrγ > 10–15Å. We therefore omit NGC 2992 from
the discussion and analysis in Sections 4 and 5.
A.2.4. NGC 1097
In NGC 1097, the first evidence for recent star formation near the nucleus was in the form of a reduction
in the stellar velocity dispersion. Emsellem et al. (2001) proposed this could be explained by the presence of
a dynamically cold nuclear disk that had recently formed stars. Direct observations of a spiral structure in
the central few arcsec, from K-band imaging (Prieto et al. 2005) and [N ii] streaming motions (Fathi et al.
2006), have since confirmed this idea. However, some issues remain open, such as why there are three spiral
arms rather than the usual two, and why gas along one of them appears to be outflowing.
Our data, at a resolution of 0.25′′ measured from the H-band non-stellar continuum, also reveal the
same spiral structure. Indeed, we find that it is traced by the morphology of the CO bandhead absorption as
well as by the 2.12µm H2 1-0 S(1) line. Interestingly, 1-0 S(1) emission is stronger where the stellar features
are weaker. This suggests that obscuration by gas and dust plays an important role. Fig. 18 shows that an
r1/4 law, typical of stellar bulges, with effective radius Reff = 0.5
′′ is a good fit to the stellar radial profile
at 0.5′′ < r < 1.8′′. It therefore seems reasonable to argue that at these radii it is only the gas that lies in a
disk. In this picture the spiral structure in the stellar continuum arises solely due to extinction of the stars
behind the disk. Extrapolating this fit, convolved with the PSF, to the nucleus indicates that at r < 0.5′′
there is at least 25% excess stellar continuum. There could be much more, given that it coincides with a
change in the dominant kinematics.
For NGC 1097 we parameterized the kinematics of the gas and stars quantitatively using kinemetry.
Based on the uniformity of the velocity field, we made the simplifying assumption that across the central 4′′
the gas lies in a single plane whose centre is coincident with the peak of the non-stellar emission. We were then
able to derive the position angle and inclination of the disk (see Section 2). The 2D kinematics of the stars
– 23 –
is traced via the CO2-0 absorption bandhead, and that of the gas through the 1-0S(1) emission line. These
independently yielded similar parameters: both gave a position angle of −49◦ and their inclinations were 43◦
and 32◦ respectively. These are fully consistent with values found by other authors (Storchi-Bergmann et al.
2003; Fathi et al. 2006). The resulting rotation curves and velocity dispersions are shown in Fig. 19. The
residuals, which can be seen in the velocity field of the gas but not the stars, and their relation to the spiral
structure described above will be discussed elsewhere (Davies et al. in prep). The important result here is
that at our spatial resolution, we find that the central stellar dispersion is σ∗ = 100 km s
−1, less than the
surrounding 150 km s−1 and also less than that in the seeing limited spectra of Emsellem et al. (2001). In
the same region we find that the rotation velocity of the gas starts to decrease rapidly, and its dispersion
increases from σgas ∼ 40 km s
−1 to ∼ 80 km s−1.
Fig. 19 also shows that while the kinematics of the stars and gas are rather different at large (> 0.5′′)
radii, they are remarkably similar at radii < 0.5′′. This certainly provides a strong indication that in the
nuclear region the stars and gas are coupled, most likely in a (perhaps thick) disk; and that the stars in this
disk, which are bright and hence presumably young, give rise to the excess stellar continuum observed.
Evidence for a recent starburst has been found by Storchi-Bergmann et al. (2005) through optical and
UV spectra. They argued that a number of features they observed could only arise from an 106 M⊙ instanta-
neous starburst, which occurred a few Myr ago and is reddened by AV = 3 mag of extinction. Using STARS
we have modeled this starburst as a 106 M⊙ burst beginning 8 Myr ago with an exponential decay timescale
of 1 Myr. The age we have used is a little older to keep the Brγ equivalent width low; and at this age, the
model predicts WBrγ = 4Å. As Fig. 20 shows, the observed Brγ is weak, although perhaps slightly resolved.
Corrected for the non-stellar continuum, we measure only WBrγ ∼ 1Å. However, the bulge population may
account for a significant fraction of the K-band stellar continuum. Correcting also for this could increase
WBrγ to 2–5Å, consistent with that of the model – assuming that the Brγ is associated with the starburst
rather than the AGN. To within a factor of a few, the scale of the model starburst is also consistent with
that measured: In the central 0.5′′ we measure a Brγ flux of 2 × 10−19 W m−2, compared to that predicted
by the model of 5× 10−19 W m−2. Given the uncertainties – factors of a few – both in the parameters of the
starburst model and also in the corrections we have applied to the data, we consider this a good agreement.
We cannot constrain the starburst further due to its compactness. Storchi-Bergmann et al. (2005) found
that it was occurring in the central 0.2′′, whereas our resolution is only 0.25′′. The Brγ emission is confined
to the central 0.4–0.5′′, although its size is hard to measure due to its weakness with respect to the stellar
absorption features. In this region the K-band stellar luminosity is 4.5× 106 L⊙. To estimate the dynamical
mass we use the mean kinematics of the stars and gas, i.e. Vrot = 40 km s
−1 (corrected for inclination) and
σ = 90 km s−1 (this is the central value, which is least biassed by bulge stars), yielding 1.4× 108 M⊙. This is
actually dominated by the black hole, which has a mass of (1.2±2)×108 M⊙ (Lewis & Eracleous 2006). The
difference between these implies a mass of gas and stars of ∼ 2× 107 M⊙, although with a large uncertainty.
The associated mass-to-light ratio is M/LK ∼ 4. On its own, this implies that over the relatively large area
that it encompasses, the maximum characteristic age for the star formation is a few hundred Myr. If one
speculates that star formation has been occurring sporadically for this timescale, then the starburst seen by
Storchi-Bergmann et al. (2005) is the most recent active episode.
In order to make a rough estimate of the supernova rate in the central region we make use of measure-
ments reported by Hummel et al. (1987). They find an unresolved component (size < 0.1′′) with 5 GHz flux
density 3.5 ± 0.3 mJy, but at lower resolution there is a 4.1 ± 0.3 mJy component of size 1′′. As discussed
in Section 3 we assume that the difference – albeit with only marginal significance – of 0.6 ± 0.4 mJy is
due to star formation in the central region, which implies a supernova rate of 6 × 10−4 yr−1 and hence
– 24 –
1010 νSN/LK ∼ 1.3, a value consistent with rather more recent star formation. Indeed, when compared to
Fig. 4, this and the low WBrγ imply a young age and short star formation timescale. For τSF = 10 Myr
the age is 60–70 Myr; for an instantaneous burst of star formation, the age would be ∼ 10 Myr, broadly
consistent with that of Storchi-Bergmann et al. (2005).
Thus, although our data do not uniquely constrain the age of the starburst in the nucleus of NGC 1097,
they do indicate that recent star formation has occurred; and they are consistent with a very young compact
starburst similar to that derived from optical and UV data.
A.2.5. NGC 1068
Evidence for a stellar core in NGC 1068 with an intrinsic size scale of ∼ 45 pc was first presented by
Thatte et al. (1997). Based on kinematics measured in large (2–4′′) apertures, they assumed the core was
virialized and estimated a mass-to-light ratio based on this assumption leading to an upper limit on the
stellar age of 1600Myr. Making a reasonable correction for an assumed old component lead to a younger
age of 500 Myr.
Stellar kinematics from optical integral field spectra (Emsellem et al. 2006a; Gerssen et al. 2006) show
evidence for a drop in the stellar velocity dispersion in the central few arcsec to σ∗ ∼ 100 km s
−1, inside a
region of higher 150–200km s−1 dispersion (presumably the bulge). Our near infrared adaptive optics data
are able to fully resolve the inner region where σ∗ drops, as shown in Fig. 21. As for NGC 1097, the velocity
distribution of the stars was derived through kinemetry, again making use of the uniformity of the stellar
velocity field to justify the simplifying assumption that the position angle and inclination do not change
significantly in the central 4′′. The derived inclination of 40◦ and position angle of 85◦ are quantitatively
similar to those found by other authors in the central few to tens of arcseconds (Emsellem et al. 2006a;
Gerssen et al. 2006; Garćıa-Lorenzo et al. 1999). The uniformity of the stellar kinematics is in contrast to
molecular gas kinematics, as traced via the 1-0 S(1) line, which are strongly perturbed and show several
distinct structures superimposed. These are too complex to permit a comparably simple analysis and will
be discussed, together with the residuals in the stellar kinematics in a future work (Mueller Sánchez et al.
in prep).
The crucial result relevant here is that at our H-band resolution of 0.10′′ we find that σ∗ reduces from
130 km s−1 at 1–2′′ to only 70 km s−1 in the very centre. That there is in the same region an excess in the
stellar continuum is demonstrated in Fig. 22. Here we show the radial profile of the stellar continuum from
both SINFONI integral field spectra out to a radius of 2′′ and NACO longslit spectra out to 5′′ (350 pc).
At radii 1–5′′, corresponding roughly to the region of high stellar dispersion measured by Emsellem et al.
(2006a), the profile is well matched by an r1/4 law, as one might expect for a bulge. At radii r < 1′′ – the
same radius at which we begin to see a discernable reduction in the stellar dispersion – the stellar continuum
increases by as much as a factor 2 above the inward extrapolation of the profile, indicating that there is extra
emission. As for NGC 1097, the combined signature of dynamically cool kinematics and excess emission is
strong evidence for a nuclear disk which has experienced recent star formation.
We can make an estimate of the characteristic age of the star formation in the central arcsec based on
the mass-to-light ratio in a similar way to Thatte et al. (1997). Because the stars appear to lie in a disk,
we estimate the dynamical mass as described in Section 3 from the stellar kinematics, using the rotation
velocity and applying a correction for the dispersion. The stellar rotation curve is essentially flat at 0.1–0.5′′,
with V∗ = 45 km s
−1 (corrected for inclination). We also take σ∗ = 70 km s
−1, which is the central value and
– 25 –
hence least biassed by the high dispersion bulge stars. These lead to a mass of 1.3× 108 M⊙ within r = 0.5
(35 pc), and a mean surface density of 3×104 M⊙ pc
−2. Correcting for the non-stellar continuum, the H-band
magnitude (which the behaviour of σ∗ indicates is dominated by the disk emission) in the same region is
11.53 mag. For H − K = 0.15 mag (Fig. 4), we find LK = 4.3 × 10
7 L⊙ and hence M/LK = 3 M⊙/L⊙.
If no star formation is on-going, this implies a characteristic age of 200–300Myr fairly independent of the
timescale (for τSF . 100 Myr, see Fig. 4) on which stars were formed. We note that this is significantly
younger than the age estimated by Thatte et al. (1997) primarily because their mass was derived using a
higher σ∗ corresponding to the bulge stars.
The assumption of no current star formation is clearly demonstrated by the Brγ map in Fig. 23. Away
from the knots of Brγ, which are associated with the coronal lines and the jet rather than possible star
formation, the equivalent width is WBrγ ∼ 4Å. This is significantly less than that for continuous star
formation of any age. Thus, while it seems likely that star formation has occurred in the last few hundred
Myr, it also seems an unavoidable conclusion that there is no current star formation.
To complete our set of diagnostics for NGC 1068, we consider also the radio continuum. This is clearly
dominated by phenomena associated with the AGN and jets, and our best estimate of the flux density away
from these features is given by the lowest contour in maps such as Figure 1 of Gallimore et al. (2004). From
this we estimate an upper limit to the 5 GHz continuum associated with star formation of 128 mJy within
r < 0.5′′. However, converting to a supernova rate and comparing to the K-band stellar luminosity yields a
limit that is not useful, being an order of magnitude above the largest expected values.
A.2.6. NGC 3783
At near infrared wavelengths, the AGN in NGC 3783 is remarkably bright. Integrated over the central
0.5′′ less than 4% of the K-band continuum is stellar. In addition, the broad Brackett lines are very strong
and dominate the H-band. Both of these phenomena are immediately clear from the H- and K-band spectra
in Fig. 24. However, it does mean that the spatial resolution can be measured easily from both the non-stellar
continuum and the broad emission lines (see Section 2). We find the K-band PSF to be symmetrical with a
FWHM of 0.17′′.
Due to the ubiquitous Brackett emission in the H-band we were unable to reliably trace the stellar
absorption features and map out the stellar continuum. Instead we have used the CO 2-0 bandhead at
2.3µm even though the dilution at the nucleus itself is extreme. The azimuthally averaged radial profile
is shown in Fig. 25 together with the PSF for reference. At radii from 0.2′′–1.6′′ (the maximum we can
measure) the profile is well fit by an r1/4 de Vaucouleurs law with Reff = 0.6
′′ (120 pc). As has been the case
previously, at smaller radii we find an excess that here is perhaps marginally resolved. Thus a substantial
fraction of the near infrared stellar continuum in the central region is likely to originate in a population of
stars distinct from the bulge.
We were unable to measure the stellar kinematics due to the limited signal-to-noise. Instead, we used
the molecular gas kinematics to estimate the dynamical mass. As before, we used kinemetry to derive the
position angle of −14◦ and the inclination in the range 35–39◦. This orientation is consistent with the
larger (20′′) scale isophotes in the J-band 2MASS image and implies that in NGC 3783 there is no significant
warp on scales of 50 pc to 4 kpc. A small inclination is also consistent with its classification as a Seyfert 1.
Adopting these values, the resulting rotation curve is shown in Fig. 26. At very small radii the rising rotation
curve may be the result of beam smearing across the nucleus. At r > 0.2′′, the falling curve suggests that
– 26 –
the rotation is dominated by the central (r < 0.2′′) mass, perhaps the supermassive black hole. We estimate
the dynamical mass within a radius of 0.3′′ (60 pc), corresponding to the point where the excess continuum
begins and also where the rotation curve appears to be unaffected by beam smearing.
Taking Vrot = 60 km s
−1 and σ = 35 km s−1 we derive a dynamical mass of Mdyn = 1.0 × 10
8 M⊙. The
black hole mass of 3 × 107 M⊙ (from reverberation mapping, Peterson et al. 2004) is only 30% of this, and
so cannot be dominating the dynamics on this scale unless its mass is underestimated. With respect to this,
we note that Peterson et al. (2004) claim the statistical uncertainty in masses derived from reverberation
mapping is about a factor 3. Alternatively, there may be a compact mass of gas and stars at r < 0.3′′.
However, including σ in the mass estimate implicitly assumes that the dispersion arises from macroscopic
motions. On the other hand, because we are observing only the hot H2, it is possible that the dispersion is
dominated by turbulence arising from shocks or UV heating of clouds that generate the 1-0 S(1) emission –
issues that are discussed in more detail by Hicks et al. (in prep.). In this case we will have overestimated
the dynamical mass. Excluding σ from the mass estimation yields Mdyn = 5 × 10
7 M⊙. We consider these
two estimates as denoting the maximum range of possible masses. Subtracting MBH then gives a mass
of stars and gas in the range (2–7)×107 M⊙, implying a mass surface density of 1700–6000M⊙ pc
−2 and
M/LK = 0.6–2.1M⊙/L⊙. Based on these ratios alone, Fig 4 indicates that the characteristic age of the star
formation may be as low as ∼ 70 Myr, although it could also be an order of magnitude greater. Without
additional diagnostics we cannot discriminate further.
We are unable to use Brγ as an additional constraint on the star formation history. Its morphology and
velocity field are similar to that of [Sivi], and rather different from the 1-0 S(1). It shows an extension to the
north which appears to be outflowing at > 50 km s−1 (Fig 27) – perhaps tracing an ionisation cone. Since
the Brγ resembles the [Sivi], it is reasonable to conclude that it too is associated with the AGN rather than
star formation. Thus the equivalent width of Brγ (with respect to the stellar continuum) of WBrγ = 30Å
represents an upper limit to that associated with star formation.
The radio continuum in the nucleus of NGC 3783 has been measured with several beam sizes at 8.5 GHz.
For a beam of 1.59′′ × 0.74′′, Morganti et al. (1999) found it was unresolved with a flux density of 8.15 ±
0.24 mJy. With a smaller ∼ 0.25′′ beam, Schmitt et al. (2001) measured a total flux density of 8.0 mJy
dominated by an unresolved component of 7.7± 0.05mJy. At smaller scales still of ∼ 0.03′′ corresponding to
6 pc, Sadler et al. (1995) placed an upper limit on the 8.5 GHz flux density of 7 mJy. Taken together, these
results imply that there is some modest 8.5 GHz radio continuum of 0.7–1 mJy extended on scales of 0.3–1′′.
Based on this we estimate a supernova rate as described in Section 3 of ∼ 0.007 yr−1, and hence a ratio
1010 νSN/LK ∼ 2. Given that the unresolved radio continuum on the smallest scales is an upper limit, the
extended component may be stronger and hence the true νSN/LK ratio may be greater than that estimated
here. Fig 4 then puts a relatively strong limit of ∼ 50 Myr on the maximum age of the star formation.
This age is fully consistent with that above associated with our lower mass estimate. The value of
WBrγ < 30Å above does not impose additional constraints, although we note that if the Brγ flux associated
with star formation is only a small fraction of the total then it would imply that the timescale over which
the star formation was active is no longer than a few times ∼ 10 Myr. Therefore in the nucleus (r < 0.3′′) of
NGC 3783 we adopt 50–70Myr as the age of the star formation and Mdyn = 2 × 10
7 M⊙ as the dynamical
mass excluding the central supermassive black hole.
– 27 –
REFERENCES
Abuter R., Schreiber J., Eisenhauer F., Ott T., Horrobin M., Gillesen S., 2006, NewAR, 50, 398
Anantharamaiah K., Viallefond F., Mohan R., Goss W., Zhao J., 2000, ApJ, 537, 613
Asari N., Vega L., Garcia-Rissmann A., González Delgado R., Storchi-Bergmann T., Cid Fernandes R., 2007,
in Galaxy Evolution Across the Hubble Time, eds Combes F., Palous J., IAU Symposium 235, 71
Baganoff F., et al., 2003, ApJ, 591, 891
Balcells M., Graham A., Domı́nguez-Palmero L., Peletier R., 2003, ApJ, 582, L79
Bender R., Burstein D., Faber S., 1992, ApJ, 399, 462
Bland-Hawthorne J., Lumsden S., Voit G., Cecil G., Weisheit J., 1997, Ap&SS, 248, 177
Bondi H., 1952, MNRAS, 112, 195
Bonnet H., et al., 2004, The ESO Messenger, 117, 17
Brand K., et al., 2006, ApJ, 663, 204
Capetti A., Axon D., Macchetto F., Marconi A., Winge C., 1999, ApJ, 516, 187
Carilli C., Wrobel J., Ulvestad J., 1998, AJ, 115, 928
Chapman S., Morris S., Alonso-Herrero A., Falcke H., 2000, MNRAS, 314, 263
Chevalier R., 1977, ARA&A, 15, 175
Chevalier R., Fransson C., 2001, ApJ, 558, L27
Cid Fernandes R., Gu Q., Melnick J., Terlevich E., Terlevich R., Kunth D., Rodrigues Lacerda R., Joguet
B., 2004, MNRAS, 355, 273
Cid Fernandes R., Heckman T., Schmitt H., González Delgado R., Storchi-Bergmann T., 2001, ApJ, 558, 81
Colina L., Alberdi A., Torrelles J., Panagia N., Wilson A., 2001, ApJ, 553, L19
Condon J., Huang Z.-P., Yin Q., Thuan T., 1991, ApJ, 378, 65
Condon J., 1992, ARA&A, 30, 575
Côté P., et al., 2006, ApJS, 165, 57
Cuadra J., Nayakshin S., Springel V., Di Matteo T., 2006, MNRAS, 366, 358
Dasyra K., et al., 2006, ApJ, 651, 835
Davies R., Sternberg A., Lehnert M., Tacconi-Garman L., 2003, ApJ, 597, 907
Davies R., Tacconi L., Genzel R., 2004a, ApJ, 602, 148
Davies R., Tacconi L., Genzel R., 2004b, ApJ, 613, 781
Davies R., Sternberg A., Lehnert M., Tacconi-Garman L., 2005, ApJ, 633, 105
– 28 –
Davies R., et al., 2006, ApJ, 646, 754
Davies R., 2007a, MNRAS, 375, 1099
Davies R., 2007b in The 2007 ESO Instrument Calibration Workshop, ESO Astrophysics Symposia,
arXiv:astro-ph/0703044
Downes D., Solomon P., 1998, ApJ, 507, 615
Eisenhauer F., et al., 2003, in Instrument Design and Performance for Optical/Infrared Ground-based Tele-
scopes, eds. Masanori I., Moorwood A., Proc. SPIE, 4841, 1548
Emsellem E., Greusard D., Combes F., Friedli D., Leon S., Pécontal E., Wozniak H., 2001, A&A, 368, 52
Emsellem E., Fathi K., Wozniak H., Ferruit P., Mundell C., Schinnerer E., 2006a, MNRAS, 365, 367
Emsellem E., 2006b, in The Central Engine of Active Galactic Nuclei, eds Ho L., Wang J.-M. (San Francisco:
Emsellem E., 2006c, in Mapping the Galaxy and Nearby Galaxies, eds Wada K., Combes F. (Springer)
Fathi K., Storchi-Bergmann T., Riffel R., Winge C., Axon D., Robinson A., Apetti A., Marconi A., 2006,
ApJ, 641, L25
Förster Schreiber N., Genzel R., Lutz D., Sternberg A., 2003, ApJ, 599, 193
Förster Schreiber N., 2000, AJ, 120, 2089
Gallimore J., Baum S., O’Dea C., 2004, ApJ, 613, 794
Garćıa-Lorenzo B., Mediavilla E., Arribas S., 1999, ApJ, 518, 190
Genzel R., Weitzel L., Tacconi-Garman L.E., Blietz M., Cameron M., Krabbe A., Lutz D., Sternberg A.,
1995, ApJ, 444, 129
Genzel R., et al., 2003, ApJ, 594, 812
Gerssen J., Allington-Smith J., Miller B., Turner J., Walker A., 2006, MNRAS, 365, 29
Glass I., 1984, MNRAS, 211, 461
González Delgado R., Heckman T., Leitherer C., 2001, ApJ, 546, 845
González Delgado R., Cid Fernandes R., 2005, in Starbursts: from 30 Doradus to Lyman Break Galaxies,
eds de Grijs R., González Delgado R., (Dordrecht: Springer), ASSL, 329, 263
Graham A., Guzmán R., 2003, AJ, 125, 2936
Granato G., de Zotti G., Silva L., Bressan A., Danese L., 2004, ApJ, 600, 580
Greenhill L., et al., 2003, ApJ, 590, 162
Gu Q., Dultzin-Hacyan D., de Diego J., 2001, RMxAA, 37, 3
Heckman T., González-Delgado R., Leitherer C., Meurer G., Krolik J., Wilson A., Koratkar A., Kinney A.,
1997, ApJ, 482, 114
http://arxiv.org/abs/astro-ph/0703044
– 29 –
Huang Z., Thuan T., Chevalier R., 1994, ApJ, 424, 114
Hummel E., van der Hulst J., Keel W., 1987, A&A, 172, 32
Ivanov V., Rieke G., Groppi C., Alonso-Herrero A., Rieke M., Engelbracht C., 2000, ApJ, 545, 190
Johuet B., Kunth D., Melnick J., Terlevich R., Terlevich E., 2001, A&A, 380, 19
Kennicutt Jr R., 1998, ApJ, 498, 541
Knapen J., 2004, in Penetrating bars through masks of cosmic dust : the Hubble tuning fork strikes a new note,
eds Block D., Puerari I., Freeman K., Groess R., Block E. (Dordrecht: Kluwer Academic Publishers),
ASSL, vol. 319, p.189
Kormendy J., Sanders D., 1992, ApJ, 390, L53
Kormendy J., Fisher D., Cornell M., Bender R., 2007, ApJ, submitted
Krajnović D., Cappellari M., de Zeeuw T., Copin Y., 2006, MNRAS, 366, 787
Krolik J., 2007, ApJ, in press
Krumholz M., McKee C., Klein R., 2006, ApJ, 638, 369
Leitherer C., et al., 1999, ApJS, 123, 3
Lenzen R., et al., 2003, in Instrument Design and Performance for Optical/Infrared Ground-based Telescopes,
eds. Masanori I., Moorwood A., Proc. SPIE, 4841, 944
Lewis K., Eracleous M., 2006, ApJ, 642, 711
Lodato G., Bertin G., 2003, A&A, 398, 517
Lonsdale C., Smith H., Lonsdale C., 1995, ApJ, 438, 632
Machalski J., Chyży K., Stawarz  L., Kozie lD., 2007, A&A, 462, 43
Maiolino R., Krabbe A., Thatte N., Genzel R., 1998, ApJ, 493, 650
Martinez-Sansigre A., Rawlings S., Lacy M., Fadda D., Jarvis M., Marleau F., Simpson C., Willott C., 2006,
MNRAS, 370, 1479
Martini P., Regan M., Mulchaey J., Pogge R., 2003, ApJ, 589, 774
Márquez I., Masegosa J., Durret F., González Delgado R., Moles M., Maza J., Pérez E., Roth M., 2003,
A&A, 409, 459
Melia F., ApJ, 387, L25
Mihos C., Hernquist L., 1994, ApJ, 437, L47
Morganti R., Tsvetanov Z., Gallimore J., Allen M., 1999, A&AS, 137, 457
Mueller Sánchez F., Davies R., Eisenhauer F., Tacconi L., Genzel R., Sternberg A., 2006, A&A, 454, 481
Nelson C., Whittle M., 1995, ApJS, 99, 67
– 30 –
Nikolaev S., Weinberg M., 1997, ApJ, 487, 885
Oliva E., Origlia L., Kotilainen J., Moorwood A., 1995, A&A, 301, 55
Origlia L., Moorwood A., Oliva E., 2003, A&A, 280, 536
Ozernoy L, Genzel R., 1996, eds Blitz L., Teuben P., (Kluwer) Proc. IAU Symposium 169, p.181
Ozernoy L., Genzel R., Usov V., 1997, MNRAS, 288, 1997
Paumard T., Maillard J.-P., Morris M., Rigaut F., 2001, A&A, 366, 466
Paumard T., et al., 2006, ApJ, 643, 1011
Peterson B., et al., 2004, ApJ, 613, 682
Pier E., Krolik J., 1992, ApJ, 399, L23
Pier E., Antonucci R., Hurt T., Kriss G., Krolik J., 1994, ApJ, 428, 124
Prieto A., Maciejewski W., Reunanen J., 2005, AJ, 130, 1472
Rigby J., Rieke G., Donley J., Alonso-Herrero A., Pérez-González P., 2006, ApJ, 645, 115
Rousset G., et al., 2003, in Adaptive Optical System Technologies II, eds. Wizinowich P., Bonaccini D., Proc.
SPIE, 4839, 140
Sadler E., Slee O., Reynolds J., Roy A., 1995, MNRAS, 276, 1273
Sarzi M., Shields J., Pogge R., Martini P., 2007, in Stellar Populations as Building Blocks of Galaxies, eds
Peletier R., Vazdekis A., IAU Symposium 241
Schmitt H., Ulvestad J., Antonuccci R., Kinney A., 2001, ApJS, 132, 199
Scoville N., et al., 2000, AJ, 119, 991
Scoville N., 2003, JKAS, 36, 167
Shapiro K., Gerssen J., van der Marel R., 2003, AJ, 126, 2707
Smith H., Lonsdale C., Lonsdale C., Diamond P., 1998, ApJ, 493, L17
Springel V., Di Matteo T., Hernquist L., 2005, MNRAS, 361, 776
Sternberg A., 1998, ApJ, 506, 721
Sternberg A., Hoffman T., Pauldrach A., 2003, ApJ, 599, 1333
Storchi-Bergmann T., González Delgado R., Schmitt H., Cid Fernandes R., Heckman T., 2001, ApJ, 559,
Storchi-Bergmann T., et al., 2003 ApJ, 598, 956
Storchi-Bergamnn T., Nemmen R., Spinelli P., Eracleous M., Wilson A., Filippenko A., Livio M., 2005, ApJ,
624, L13
Thatte N., Quirrenbach A., Genzel R., Maiolino R., Tecza M., 1997, ApJ, 490, 238
– 31 –
Thompson T., Quataert E., Murray N., 2005, ApJ, 630, 167
Vazquez G., Leitherer C., 2005, ApJ, 621, 695
Wehrle A., Morris M., 1988, AJ, 95, 1689
Winters J., Le Bertre T., Jeong K., Nyman L.-Å., Epchtein N., 2003, A&A, 409, 715
Woo J.-H., Urry C., 2002, ApJ, 579, 530
Wozniak H., Combes F., Emsellem E., Friedli D., 2003, A&A, 409, 469
This preprint was prepared with the AAS LATEX macros v5.2.
– 32 –
Table 1. Table of Observations
Object Banda Res.b (′′) Date Instrument
Mkn 231c H 0.176 May ’02 Keck, NIRC2
NGC 7469c K 0.085 Nov ’02 Keck, NIRSPAO
K 0.15 Jul ’04 VLT, SINFONI
Circinusc K 0.22 Jul ’04 VLT, SINFONI
NGC 3227c K 0.085 Dec ’04 VLT, SINFONI
IRAS 05189-2524 H 0.12 Dec ’02 VLT, NACO
NGC 2992 K 0.30 Mar ’05 VLT, SINFONI
NGC 1097 H 0.245 Oct ’05 VLT, SINFONI
NGC 1068 H 0.10 Oct ’05 VLT, SINFONI
H 0.13 Dec ’02 VLT, NACO
NGC 3783 K 0.17 Mar ’05 VLT, SINFONI
aBand used for determining the quantitative star formation proper-
ties. NGC 1097, NGC 1068, and NGC 3783 were actually observed in H
and K bands.
bSpatial resolution (FWHM) estimated from the data itself, using the
methods described in Section 2.
cReferences to detailed studies of individual objects: Mkn 231
(Davies et al. 2004b), NGC 7469 (Davies et al. 2004a), Circinus,
(Mueller Sánchez et al. 2006), NGC 3227 (Davies et al. 2006).
– 33 –
Table 2. Summary of basic data for AGN
Object Classificationa Distance logLbol
b logMBH
Ref.c
Mpc for MBH
Mkn 231 ULIRG, Sy 1, QSO 170 12.5 7.2 1
NGC 7469 Sy 1 66 11.5 7.0 2
Circinus Sy 2 4 10.2 6.2 3
NGC 3227 Sy 1 17 10.2 7.3 4
IRAS 05189-2524 ULIRG, Sy 1 170 12.1 7.5 1
NGC 2992 Sy 1 33 10.7 7.7 5
NGC 1097 LINER, Sy 1 18 10.9 8.1 6
NGC 1068 Sy 2 14 11.5 6.9 7
NGC 3783 Sy 1 42 10.8 7.5 2
a Classifications are taken primarily from the NASA/IPAC Extragalactic Database. In
addition, we have labelled as Seyfert 1 those for which we have observed broad (i.e. FWHM
> 1000 km s−1) Brγ; see also Fig 1.
b Calculated in the range 8–1000µm from the IRAS 12–100µm flux densities; with an
additional correction for optical and near-infrared luminosity in cases where appropriate.
c References for black hole masses: (1) Dasyra et al. (2006); (2) Peterson et al.
(2004); (3) Greenhill et al. (2003); (4) Davies et al. (2006); (5) Woo & Urry (2002); (6)
Lewis & Eracleous (2006); (7) Lodato & Bertin (2003)
Table 3. Measured & Derived Properties of the Nucleia
Object radius logLK∗
Σdyn WBrγ Mdyn/LK 10
10νSN/LK age 〈SFR〉
′′ pc 104M⊙ pc
−2 Å M⊙/L⊙ yr
−1L−1
Myr M⊙ yr
−1 kpc−2
Mkn 231b 0.6 480 9.3 9.8 0.9 — 3.1 20 120–250 25–50
NGC 7469 0.4 128 8.5 8.7 1.0 11 1.6 3 110–190 50–100
Circinus 0.4 8 6.2 7.5 17 30 23 1.5 80 ∼ 70
NGC 3227c 0.4 32 7.8 8.0 3.7 4 1.9 2.2 40 ∼ 380
IRAS 05189-2524 0.55 450 9.3 — — 4 — 5 50–100 30–70
NGC 2992d 0.4 64 7.5 — — <12 — 1 — —
NGC 1097e 0.25 22 6.7 8.2 1.3 1 4.5 1.4 8 ∼ 80
NGC 1068 0.5 35 7.6 8.1 3.4 4 3.0 <20 200–300 90–170
NGC 3783f 0.3 60 7.5 7.3 0.2 <30 0.6 2 50–70 30–60
a The methods used to measure these quantities (within the radii given) are described in Section 3. Specific issues associated with
individual objects are discussed in Appendix A.
b Mdyn depends strongly on even small changes to the inclination; here it is given for i = 10
◦. Correcting Mdyn for an estimate of the gas
mass given in Downes & Solomon (1998) yields M/LK = 2.3M⊙/L⊙.
c The best star formation models indicate that M/LK is much less than the limit given here using the dynamical mass.
d It is likely that much of the narrow Brγ in the nuclear region here is associated with an ionisation cone. In addition, the high stellar
velocity dispersion, even on the smallest scales we have been able to measure, suggests that the K-band light is dominated by the bulge.
e Correcting WBrγ for the old stellar population would probably yield a value in the range 2–5Å. Even on this scale the dynamical mass
is dominated by the supermassive black hole. Both Σdyn and M/LK are estimated after subtracting MBH.
f Much of the Brγ here is outflowing and hence associated with the AGN. Mdyn is derived from gas kinematics as described in the text.
Both Σdyn and M/LK are estimated after subtracting MBH.
– 35 –
Fig. 1.— K-band spectra showing broad Brγ emission in 3 AGN which are not usually classified as Seyfert 1.
Top: NGC 3227 (0.25′′ aperture); Middle: IRAS 05189−2524 (1′′ aperture); Bottom: NGC 2992 (0.5′′ aper-
ture). The most prominent emission and absorption features are marked.
Fig. 2.— Left: equivalent width of the CO 2-0 and CO 6-3 features for various stars and galaxies. The late-
type supergiant stars (skeletal star shapes) and giant stars (open star shapes) are taken from Origlia et al.
(1993). The galaxies (denoted ‘E’ for elliptical, ‘S’ for spiral, and ‘H’ for star forming Hii galaxy) are from
Oliva et al. (1995). The dashed box encloses the region for which there is no more than 20% deviation
from each of the values WCO2−0 = 12Å and WCO6−3 = 4.5Å. Centre and right: calculated WCO6−3 and
WCO2−0 respectively from STARS for several different star formation histories. Each line is truncated when
the cluster luminosity falls below 1/15 of its maximum. In each case, the dashed lines show typical values
adopted, and the dotted lines denote a range of ±20%.
– 36 –
Fig. 3.— Left: H-K colour of star clusters with different star formation timescales and ages, as calculated by
STARS. Right: Ratio of bolometric to K-band luminosity. Although the range of 20–200 initially appears
large, the uncertainty on an intermediate value of 60 is only 0.3 dex. This is small compared to the range
of interest in the paper, which is several orders of magnitude.
Fig. 4.— Various diagnostics calcuated with STARS for several star formation timescales, as functions of
age: Brγ equivalent width, supernova rate, and mass-to-light ratio. Note that all are normalised to the
K-band stellar continuum; and that LK is the total luminosity in the 1.9–2.5µm band in units of bolometric
solar luminosity (1Lbol = 3.8 × 10
26 W ), rather than the other frequently used monochromatic definition
which has units of the solar K-band luminosity density.
– 37 –
Fig. 5.— Size scales of nuclear star forming regions. The profile has been determined from the CO absorption
features in the H or K band, which are approximately independent of star formation history (see text for
details). For longslit data, spatial profiles have been averaged; for integral field data, azimuthally averaged
proifles are shown. For NGC1068, data at two different pixel scales are shown (corresponding to the solid
and dashed brown lines). The horizontal dotted line is drawn at half-maximum height, to assist in estimating
size scales by eye.
– 38 –
Fig. 6.— A STARS star formation model based on the main characteristics of the observed starbursts,
which is illustrative of a ‘typical’ nuclear starburst that we have observed. The scaling is fixed as 2× 109 M⊙
(typical of that within a 30 pc radius; Fig. 9) at an age of 100 Myr (the typical age in Table 3). The star
formation timescale is τSF = 10 Myr to reproduce a low Brγ equivalent width. The panels are, from top left:
(a) star formation rate (SFR); (b) number of ionising photons (QLyc, proportional to the Brγ luminosity);
(c) supernova rate (SNR); (d) bolometric luminosity (Lbol); (e) K-band luminosity (LK); (f) cumulative
number of supernovae; (g) cumulative mass that has been recycled back into the ISM by supernovae and
winds; (h) mass loss rate from stars; (i) rate at which the lost mass can in principle be accreted onto a central
supermassive black hole (due to its outflow speed; see text for details). In the last two panels, the mass loss
is split into that due to OB and Wolf-Rayet stars and supernovae (dotted lines), and that due to late-type
and AGB stars (dashed lines). Stellar mass loss in STARS is accounted for at the end of each star’s life as
the difference in mass between the original (ZAMS) stellar mass and the remnant mass as the end-product
of its stellar evolution, as described in Sternberg (1998). For this reason, the mass loss rates from OB and
Wolf-Rayet stars do not appear explicitly in panel (h) at very young cluster ages.
– 39 –
Fig. 7.— Mean enclosed mass surface density as a function of radius. The points are from data given
in Table 3; the dashed lines represent the mass models derived for NGC 7469 and Mkn 231 (Davies et al.
2004a,b). The galaxies all follow the same trend towards increasing densities in the central regions.
Fig. 8.— Integrated bolometric luminosity of the young stars Lbol∗ as a fraction of that of the galaxy
Lbol−galaxy, plotted as a function of radius. Lbol∗ is calculated from the stellar LK or LH assuming that,
on the generally small scales here, all the near infrared stellar continuum originates in the young stars. On
10 pc scales the contribution of young stars is at most a few percent of the galaxy’s total luminosity, while
on kpc scales it may be significant and hence comparable to the AGN luminosity. For NGC 1068, data at
two different pixel scales are shown (corresponding to the solid and dashed brown lines).
– 40 –
Fig. 9.— Integrated bolometric luminosity of young stars Lbol∗ (see Fig. 8) as a function of radius. For
comparison, the dotted line has constant surface brightness. For NGC 1068, data at two different pixel scales
are shown (corresponding to the solid and dashed brown lines).
Fig. 10.— Luminosity to mass ratio calculated by STARS as a function of age for different star formation
timescales. The dotted lines show how the ratio would vary if gas were fed in to a cluster at the same rate
as it was converted into stars. The solid lines assume that the gas is present at the start, but at the end has
all been processed in stars. A cluster can only exceed 500 L⊙/M⊙ for a timescale of ∼ 10 Myr.
– 41 –
Fig. 11.— Graph showing how the luminosity of an AGN might be related to the age of the most recent
episode of nuclear star formation. On the left is shown the luminosity in solar units; on the right, it is with
respect to the Eddington luminosity for the black hole. Generally the luminosity of the AGN is not well
known and so we have approximated it by 0.5Lbol, and adopted an uncertainty of a factor 2. The starburst
age refers to our best estimate of the most recent episode of star formation within the central 10–100pc, as
given in Table 3. See the text for details of the adopted uncertainties.
– 42 –
Fig. 12.— Spectrum (in grey) of IRAS 05189-2524, integrated over two 0.22′′ wide sections centered ±0.27′′
either side of the nucleus, and which have been shifted to match their velocities. Overplotted in black is
a fit to the continuum, comprising spectra of various supergiant stars. The main absorption and emission
features have been identified.
Fig. 13.— Spatial profiles of non-stellar (left), stellar (centre), and total (right) continuum for IRAS 05189-
2524 (1′′ = 800 pc). The first two have been derived at each point along the spatial extent of the slit from
the spectral slope and the stellar absorption features respectively. A comparison of their sum (dashed line
in right panel) to the total continuum indicates that the decomposition appears to be reasonable.
– 43 –
Fig. 14.— Maps of NGC 2992 (1′′ = 160 pc). From left to right: stellar continuum, narrow Brγ, Brγ velocity
(−150 to +150 km s−1), and Brγ dispersion (0 to 200 km s−1). For reference, on each panel are superimposed
contours of the Brγ flux (8, 16, 32, and 64% of the peak). The symbol plotted on the map of the line flux
indicates the centre of the broad Brγ and non-stellar emission. The narrow Brγ emission extends far more
to the north west than the stellar continuum. And, particularly on the south western edge, it exhibits a
blue shifted velocity and high dispersion. All these are consistent with an interpretation as the apex of an
ionisation cone.
Fig. 15.— Radial profiles of velocity and dispersion for the stars in NGC 2992 (1′′ = 160 pc). The 2D maps
were analysed using the kinemetric technique described by Krajnović et al. (2006) which yielded a position
angle of 24◦ and an inclination of ∼ 40◦, not dissimilar to the isophotal values of 30◦ and 50◦ respectively.
The rotation curve has been corrected for the inclination.
– 44 –
Fig. 16.— Radial profile of the stellar continuum in NGC 2992 (1′′ = 160 pc), derived from isophotal analysis.
Solid circles denote the stellar continuum (i.e. already corrected for the non-stellar component). Overplotted
with triangles are an r1/4 law (top panel) and an exponential profile (bottom panel). The profiles were fitted
at radii r > 0.5′′ and extrapolated inwards, convolved with the PSF which is shown as open squares. Both
fits are equally good at r > 0.5′′, but only the exponential suggests there might be excess continuum at the
centre, arising from a distinct stellar population. This is therefore inconclusive.
– 45 –
Fig. 17.— Spectrum of NGC 1097 (thick grey line), extracted within an aperture of radius 0.2′′ and scaled
arbitrarily. Overdrawn (thin black line) is a match to the stellar continuum constructed from template
spectra of several late-type supergiant stars and a blackbody function representing the non-stellar component.
Notably, Brγ in the nucleus is extremely weak even in the nucleus.
Fig. 18.— Radial profile of the stellar continuum in NGC 1097 (1′′ = 80 pc). The solid circles denote the
stellar continuum (i.e. already corrected for the non-stellar component). The triangles denote an r1/4 profile
fitted to radii r > 0.5′′ and extrapolated inwards. This model has been convolved with the PSF, shown as
open squares for comparison. Note that even though an exponential profile might match the data equally
well, an r1/4 profile provides a stronger constraint on whether there is excess continuum at the centre.
– 46 –
Fig. 19.— Radial profiles of velocity and dispersion for the gas and stars in NGC 1097 (1′′ = 80 pc). The 2D
maps were created by convolving template spectra (i.e unresolved line profile for the gas, stellar template
for the stars) with a Gaussian and minimising the difference with respect to the galaxy spectrum at each
spatial pixel. These were then analysed using the kinemetric technique described by Krajnović et al. (2006)
which yielded the same position angle of −49◦ for the gas and stars, and similar inclinations of 32◦ and 43◦
respectively. The rotation curve has been corrected for the inclination.
Fig. 20.— Maps of K-band non-stellar continuum (left), stellar continuum (centre), and Brγ line flux (right)
for the central few arcsec of NGC 1097 (1′′ = 80 pc). In each case, the centre (as defined by the non-stellar
continuum) is marked by a crossed circle. The colour scale is shown on the right, as percentage of the peak
in each map.
– 47 –
Fig. 21.— Radial profiles of velocity and dispersion for stars in NGC 1068 (1′′ = 70 pc). The 2D maps were
then analysed using kinemetry Krajnović et al. (2006), yielding an inclination of 40◦ and a position angle of
85◦. The rotation curve has been corrected for the inclination. Also plotted for comparison are azimuthally
averaged radial profiles of the H-band stellar luminosity and the PSF.
Fig. 22.— Radial profiles of the H-band stellar luminosity in NGC 1068 (1′′= 70 pc) from NACO longslit
data at position angle 0◦ and 90◦ (green and red symbols; squares and triangles deonte opposite sides of the
nucleus). To trace the profile to larger radius, open symbols are each the mean of 9 points. Superimposed
are SINFONI data from Fig. 21 (blue crosses, flux scaled to match). The dashed line denotes an r1/4 law
with Reff = 1.5
′′ to match the outer profile. At r < 1′′, the stellar continuum reveals an excess above the
inward extrapolation of this profile.
– 48 –
Fig. 23.— Maps of the central few arcsec of NGC 1068 (1′′ = 70 pc): H-band non-stellar continuum (far left)
and stellar continuum (centre left); also Brγ line flux (center right) and Brγ equivalent width (far right).
In each case, the centre (as defined by the non-stellar continuum) is marked by a crossed circle. The colour
scale is shown on the right, as percentage of the peak in each map (and also as WBrγ in Å).
– 49 –
Fig. 24.— H- and K-band spectra of the central 1′′ of NGC 3783, with the prominent emission and absorption
features labelled. In the H-band it is challenging to measure the stellar absorption due to the very strong
brackett emission from the AGN’s broad line region. Instead we have used the K-band CO 2-0 bandhead
even though the dilution at this wavelength is extreme.
– 50 –
Fig. 25.— Radial profile of the stellar continuum in NGC 3783 (1′′ = 200 pc). The solid circles denote the
stellar continuum (i.e. already corrected for the non-stellar component). The triangles denote an r1/4 profile
fitted to radii 0.2 < r < 1.6′′ and extrapolated inwards. This model has been convolved with the PSF, shown
as open squares for comparison. Note that even though an exponential profile might match the data equally
well, an r1/4 profile provides a stronger constraint on whether there is excess continuum at the centre.
Fig. 26.— Rotation curve derived from the the H2 1-0 S(1) velocity field in NGC 3783 (1
′′ = 200 pc).
Also shown is the dispersion as a function of radius. The velocity field was analysed using kinemetry
(Krajnović et al. 2006) which yielded a major axis of about −14◦ and an inclination in the range 35–39◦.
The drop in velocity at r < 0.15′′ maybe due to beam smearing across the nucleus.
– 51 –
Fig. 27.— Images of the central 2′′ of NGC 3783 (1′′ = 200 pc). Top row from left: non-stellar continuum,
H2 1-0 S(1) line flux, H2 1-0 S(1) velocity. Bottom row from left: narrow Brγ line flux, [Sivi] line flux, [Sivi]
velocity. The Brγ velocity field is similar to that of [Sivi] and shows an outflow of >50 km s−1 to the north.
This is in contrast to the 1-0 S(1) velocity field which traces rotation.
	Introduction
	Sample, Observations, Data Processing
	Sample Selection
	Observations & Reduction
	PSF Estimation
	Emission/Absorption Line Characterisation
	Quantifying the Star Formation
	Isolating the stellar continuum
	Stellar colour and luminosity
	Specific Star Formation Diagnostics
	Br equivalent width
	Supernova rate
	Mass-to-light ratio
	Properties of Nuclear Star Formation
	Starburst-AGN Connection
	Conclusions
	Individual Objects
	Summary of Star Forming Properties of Galaxies already Studied
	Mkn 231
	Circinus
	NGC 3227
	Star Forming Properties of Galaxies with New Data
	NGC 7469
	IRAS 05189-2524
	NGC 2992
	NGC 1097
	NGC 1068
	NGC 3783

  We analyse star formation in the nuclei of 9 Seyfert galaxies at spatial
resolutions down to 0.085arcsec, corresponding to length scales of less than
10pc in some objects. Our data were taken mostly with the near infrared
adaptive optics integral field spectrograph SINFONI. The stellar light profiles
typically have size scales of a few tens of parsecs. In two cases there is
unambiguous kinematic evidence for stellar disks on these scales. In the
nuclear regions there appear to have been recent - but no longer active -
starbursts in the last 10-300Myr. The stellar luminosity is less than a few
percent of the AGN in the central 10pc, whereas on kiloparsec scales the
luminosities are comparable. The surface stellar luminosity density follows a
similar trend in all the objects, increasing steadily at smaller radii up to
10^{13}L_sun/kpc^2 in the central few parsecs, where the mass surface density
exceeds 10^4M_sun/pc^2. The intense starbursts were probably Eddington limited
and hence inevitably short-lived, implying that the starbursts occur in
multiple short bursts. The data hint at a delay of 50--100Myr between the onset
of star formation and subsequent fuelling of the black hole. We discuss whether
this may be a consequence of the role that stellar ejecta could play in
fuelling the black hole. While a significant mass is ejected by OB winds and
supernovae, their high velocity means that very little of it can be accreted.
On the other hand winds from AGB stars ultimately dominate the total mass loss,
and they can also be accreted very efficiently because of their slow speeds.

Introduction
During recent years there has been increasing evidence for a connection between active galactic nuclei
(AGN) and star formation in the vicinity of the central black holes. This subject forms the central topic of
this paper, and is discussed in Sections 4 and 5.
A large number of studies have addressed the issue of star formation around AGN. Those which have
probed closest to the nucleus, typically on scales of a few hundred parsecs, have tended to focus on Seyferts
– notably Seyfert 2 galaxies – since these are the closest examples (Sarzi et al. 2007; Asari et el. 2007;
1Based on observations at the European Southern Observatory VLT (60.A-9235, 070.B-0649, 070.B-0664, 074.B-9012, 076.B-
0098).
http://arxiv.org/abs/0704.1374v2
– 2 –
González Delgado, et al. 2005; Cid Fernandes et al. 2004; González Delgado, et al. 2001; Gu et al. 2001;
Joguet et al. 2001; Storchi-Bergmann et al. 2001; Ivanov et al. 2000). The overall conclusion of these studies
is that in 30–50% cases the AGN is associated with young (i.e. age less than a few 100 Myr) star formation.
While this certainly implies a link, it does not necessarily imply any causal link between the two phenomena.
Instead, it could more simply be a natural consequence of the fact that both AGN and starburst require gas
to fuel them. And that in some galaxies this gas has fallen towards the nucleus, either due to an interaction
or secular evolution such as bar driven inflow.
One aspect which must be borne in mind when interpreting such results, and which has been pointed
out by Knapen (2004), is the discrepancy in the scales involved. AGN and starburst phemonena occur on
different temporal and spatial scales; and observations are sensitive to scales that are different again. For
example, star formation has typically been studied on scales of several kiloparsecs down to a few hundred
parsecs. In contrast, accretion of gas onto an AGN will occur on scales much less than 1 pc. Similarly, the
shortest star formation timescales that most observations are sensitive to are of order 100 Myr to 1 Gyr. On
the other hand, in this paper we show that the active phase of star formation close around a black hole is
typically rather less than 100 Myr. Correspondingly short accretion timescales for black holes are reflected
in the ages of jets which, for a sample of radio galaxies measured by Machalski et al. (2007), span a range
from a few to 100 Myr. In Seyfert galaxies the timescales are even shorter, as typified by NGC 1068 for which
Capetti et al. (1999) estimate the age of the jets to be only ∼ 0.1 Myr. That the putative causal connection
between AGN and starbursts might occur on relatively small spatial scales and short timescales can help us
to understand why no correlation has been found between AGN and (circum-)nuclear starbursts in general.
It is simply that the circumnuclear activity on scales greater than a few hundred parsecs is, in most cases,
too far from the AGN to influence it, or be strongly influenced by it (cf Heckman et al. 1997).
In this paper we redress this imbalance. While the optical spectroscopy pursued by many authors
allows a detailed fitting of templates and models to the stellar features, we also make use of established
star formation diagnostics and interpret them using starburst population synthesis models. Observing at
near infrared wavelengths has brought two important advantages. The optical depth is 10 times less than at
optical wavelengths, and thus our data are less prone to the effects of extinction which can be significant in
AGN. And we have employed adaptive optics to reach spatial resolutions of 0.1–0.2′′, bringing us closer to the
nucleus. Applying these techniques, we have already analysed the properties of the nuclear star formation in
a few objects (Davies et al. 2004a,b, 2006; Mueller Sánchez et al. 2006). Here we bring those data together
with new data on 5 additional objects. Our sample enables us to probe star formation in AGN from radii
of 1 kpc down to less than 10 pc. Our aim is to ascertain whether there is evidence for star formation on
the smallest scales we can reach; and if so, to constrain its star formation history. Ultimately, we look at
whether there are indications that the nuclear starburst and AGN are mutually influencing each other.
In §2 we describe the sample selection, observations, data reduction, PSF estimation, and extraction of
the emission and absorption line morphologies and kinematics. In §3 we discuss the observational diagnostics
and modelling tools. Brief analyses of the relevant facets of our new data for the individual objects are
provided in Appendix A, where we also summarise results of our previously published data, re-assessing
them where necessary to ensure that all objects are analysed in a consistent manner. The primary aims of
our paper are addressed in §4 and §5. In §4 we discuss global results concerning the existence and recent
history of nuclear star formation for our whole sample. In §5 we discuss the implications of nuclear starbursts
on the starburst-AGN connection. Finally, we present our conclusions in §6.
– 3 –
2. Sample, Observations, Data Processing
2.1. Sample Selection
The AGN discussed in this paper form a rather heterogeneous group. They include type 1 and type 2
Seyferts, ULIRGs, and even a QSO, and do not constitute a complete sample. In order to maximise the size
of the sample, we have combined objects on which we have already published adaptive optics near infrared
spectra with new observations of additional targets.
Source selection was driven largely by technical considerations for the adaptive optics (AO) system,
namely having a nucleus bright and compact enough to allow a good AO correction. This is actually a
strength since it means that 7 of the 9 AGN are in fact broad line objects – as given either by the standard
type 1 classification or because there is clear broad (FWHM > 1000km s−1) Brγ emission in our spectra.
Fig. 1 shows broad Brγ in K-band spectra of 3 AGN that are not usually classified as broad line galaxies.
This is in contrast to most other samples of AGN for which star formation has been studied in detail, and
avoids any bias that might arise from selecting only type 2 Syeferts. That there may be a bias arises from the
increasing evidence that the obscuration in perhaps half of type 2 AGN lies at kpc scales rather than in the
nucleus, which may be caused by spatially extended star formation in the galaxy disk (Brand et al. 2007;
Martinez-Sansigre et al. 2006; Rigby et al. 2006). Such AGN do not fit easily into the standard unification
scheme (and perhaps should not really be considered type 2 objects). Because broad lines can be seen in
the infrared, we know that we are seeing down to the nuclear region and hence our results are not subject
to any effects that this might otherwise introduce.
It is exactly broad line AGN for which little is known about the nuclear star formation, because the glare
of the AGN swamps any surrounding stellar light in the central arcsec. As a result, most studies addressing
star formation close to AGN have focussed on type 2 Seyferts. Adaptive optics makes it possible to confine
much of the AGN’s light into a very compact region, and to resolve the stellar continuum around it. The
use of adaptive optics does give rise to one difficulty when attempting to quantify the results in a uniform
way, due to the different resolutions achieved – which is a combination of both the distance to each object
(i.e. target selection) and the AO performance. As a result, the standard deviation around the logarithmic
mean resolution of our sample (excluding NGC 2992, see Section A) of 22 pc is a factor of 3. However, this
has enabled us to study the centers of AGN across nearly 3 orders of magnitude in spatial scale, from 1 kpc
in the more distant objects to only a few parsecs in the nearby objects with the best AO correction.
2.2. Observations & Reduction
A summary of our observations is given in Table 1. A description of observations and processing of the
new data is given below.
Data for IRAS 05189-2524 and NGC 1068 were taken in December 2002 at the VLT with NACO, an
adaptive optics near infrared camera and long slit spectrograph (Lenzen et al. 2003; Rousset et al. 2003).
Since IRAS 05189-2524 is nearly face on, there is no strongly preferred axis and the slit was oriented north-
south; for NGC 1068 two orientations were used, north-south and east-west. In all cases the slit width was
0.086′′, yielding a nominal resolution of R ∼ 1500 with the wide-field camera (pixel scale 0.054′′) and medium
resolution grism. The galaxy was nodded back and forth along the slit by 10′′ to allow sky subtraction. For
IRAS 05189-2524, 12 integrations of 300 sec were made; for NGC 1068 12 integrations of 300 sec were made
at one position angle, and 14 frames of 200 sec at the other. All data were reduced and combined, using
– 4 –
standard longslit techniques in IRAF, to make the final H-band spectrum.
Data for NGC 7469, NGC 2992, NGC 1097, NGC 1068, and NGC 3783 were taken during 2004–2005 at
the VLT with SINFONI, an adaptive optics near infrared integral field spectrograph (Eisenhauer et al. 2003;
Bonnet et al. 2004). Data were taken with various gratings covering the H and K bands either separately
(R ∼ 4000) or together (R ∼ 1500). The pixel scales were 0.125′′×0.25′′ or 0.05′′×0.1′′, depending on the
trade-offs between field of view, spatial resolution, and signal-to-noise ratio. Individual exposure times are in
the range 50–300 sec depending on the object brightness. Object frames were interspersed with sky frames,
usually using the sequence O-S-O-O-S-O, to facilitate background subtraction. The data were processed
using the dedicated spred software package (Abuter et al. 2006), which provides similar processing to that
for longslit data but with the added ability to reconstruct the datacube. The data processing steps are
as follows. The object frames are pre-processed by subtracting sky frames, flatfielding, and correcting bad
pixels (which are identified from dark frames and the flatfield). The wavemap is generated, and edges and
curvature of the slitlets are traced, all from the arclamp frame. The arclamp frame is then reconstructed
into a cube, which is checked to ensure that the calibration is good. The pre-processed object frames are
then also reconstructed into cubes, spatially shifted to align them using the bright nucleus as a reference,
and combined. In some cases the final cube was spatially smoothed using a 3 × 3 median filter. Estimation
of the spatial resolution (see below) was always performed after this stage.
In some cases, the strong near-infrared OH lines did not subtract well. With longer exposure times this
is to be expected since the timescale for variation of the OH is only 1–2 mins. If visual inspection of the
reconstructed cubes showed signs of over- or under-subtraction of the OH lines, these cubes were reprocessed
using the method described in Davies (2007a).
Standard star frames are similarly reconstructed into cubes. Telluric correction and flux calibration were
performed using B stars (K-band) or G2V stars (H-band). In addition, flux calibration was cross-checked in
3′′ apertures using 2MASS data, and in smaller 1–3′′ apertures using broad-band imaging from NACO or
HST NICMOS. Agreement between cubes with different pixel scales, and also with the external data, was
consistent to typically 20%.
2.3. PSF Estimation
There are a multitude of ways to derive the point spread function (PSF) from adaptive optics data,
five of which are described in Davies (2007b). With AGN, it is usually possible to estimate the PSF from
the science data itself, removing any uncertainty about spatial and temporal variations of the PSF due to
atmospheric effects. Typically one or both of the following methods are employed on the new data presented
here. If a broad emission line is detected, this will always yield a measure of the PSF since the BLR of
Seyfert galaxies has a diameter that can be measured in light days. Alternatively, the non-stellar continuum
will provide a sufficiently good approximation in all but the nearest AGN since at near infrared wavelengths
it is expected to originate from a region no more than 1–2 pc across.
In every case we have fit an analytical function to the PSF. Since the Strehl ratio achieved is relatively
low, even a Gaussian is a good representation. We have used a Moffat function, which achieves a better fit
because it also matches the rather broad wings that are a characteristic of partial adaptive optics correction.
The PSF measured for NGC 3227, which is shown in Fig. 1 of Davies et al. (2006), can be considered typical.
If one applies the concept of ‘core plus halo’ to this PSF, then the Gaussian fit would represent just the core
while the Moffat fit the entire ‘core plus halo’. Integrating both of these functions indicates that about 75%
– 5 –
of the flux is within the ‘core’, and it is thus this component which dominates the PSF. In this paper, a more
exact representation of the PSF is not needed since we have not performed a detailed kinematic analysis,
and we have simply used the Moffat to derive a FWHM for the spatial resolution. The resolutions achieved
are listed in Table 1.
2.4. Emission/Absorption Line Characterisation
The 2D distribution of emission and absorption features has been found by fitting a function to the
continuum-subtracted spectral profile at each spatial position in the datacube. The function was a convolu-
tion of a Gaussian with a spectrally unresolved template profile – in the case of emission lines it was an OH
sky emission line, and for stellar absorption features we made use of template stars observed in the same con-
figuration (pixel scale and grism). A minimisation was performed in which the parameters of the Gaussian
were adjusted until the convolved profile best matched the data. During the minimisation, pixels in the data
that consistently deviated strongly from the data were rejected. The uncertainties were boot-strapped using
Monte Carlo techniques, assuming that the noise is uncorrelated and the intrinsic profile is well represented
by a Gaussian. The method involves adding a Gaussian with the derived properties to a spectral segment
that exhibits the same noise statistics as the data, and refitting the result to yield a new set of Gaussian
parameters. After repeating this 100 times, the standard deviation of the center and dispersion were used
as the uncertainites for the velocity and line width.
The kinematics were further processed using kinemetry (Krajnović et al. 2006). This is a parameteri-
sation (i.e. a mathematical rather than a physical model) of the 2D field. As such, beam smearing is not
a relevant issue to kinemetry, which yields an analytical expression for the observed data. Of course, when
the coefficients of this expression are interpreted or used to constrain a physical model, then beam smearing
should be considered. Mathematically, the kinemetry procedure fits the data with a linear sum of sines and
cosines with various angular scalings around ellipses at each radius. We have used it for 3 purposes: to de-
termine the best position angle and axis ratio for the velocity field, to remove high order noise from the raw
kinematic extraction, and to recover the velocity and dispersion radial profiles. In all of the cases considered
here, the kinematic centre of the velocity field was assumed to be coincident with the peak of the non-stellar
continuum. In addition, the uniformity of the velocity field permitted us to make the simplifying assumption
of a single position angle and axis ratio – i.e. there is no evidence for warps or twisted velocity contours.
We then derived the position angle and inclination of the disk by minimising the A1 and B3 parameters
respectively (see Krajnović et al. 2006 for a description of these). The rotation curves were recovered by
correcting the measured velocity profile for inclination. We have assumed throughout the paper that the
dispersion is isotropic, and hence no inclination correction was applied to the dispersion that was measured.
The innermost parts of the kinematics derived as above are of course still affected by beam smearing.
In general, the central dispersion cannot necessarily be taken at face value since it may either be artificially
increased by any component of rotation included within the beam size, or decreased if neighbouring regions
within the beam have a lower dispersion. In the galaxies we have studied, there are two aspects which
mitigate this uncertainty: the rotation speed in the central region is much less than the dispersion and so
will not significantly alter it; and when estimating the central value we consider the trend of the dispersion
from large radii, where the effect of the beam is small, to the center. For the basic analyses performed here,
we have therefore adopted the central dispersion at face value. More detailed physical models for the nuclear
disks, which properly account for the effects of beam smearing, will be presented in future publications.
Lastly, we emphasize that the impact of the finite beam size on the derived rotation curve does not affect
– 6 –
our measurement of the dynamical mass. The reason is that, for all the dynamical mass estimates we make,
the mass is estimated at a radius much large than the FWHM of the PSF – as can be seen in the relevant
figures.
3. Quantifying the Star Formation
In this section we describe the tools of the trade used to analyse the data, and which lead us to the
global results presented in Section 4. Specific details and analyses for individual objects can be found in
Appendix A. We use the same methods and tools for all the objects to ensure that all the data are analysed
in a consistent manner.
Perhaps the most important issue is how to isolate the stellar continuum, which is itself a powerful
diagnostic. In addition, we use three standard and independent diagnostics to quantify the star formation
history and intensity in the nuclei of these AGN. These are the Brγ equivalent width, supernova rate, and
mass-to-light ratio. Much of the discussion concerns how we take into account the contribution of the AGN
when quantifying these parameters. We also consider what impact an incorrect compensation could have on
interpretation of the diagnostics.
We model these observational diagnostics using the stellar population and spectral synthesis code STARS
(e.g. Sternberg 1998; Sternberg et al. 2003; Förster Schreiber et al. 2003; Davies et al. 2003, 2005). This code
calculates the distribution of stars in the Hertzsprung-Russell diagram as a function of age for an assumed
star formation history. We usually assume an exponentially decaying star formation rate, which has an
associated timescale τSF. Spectral properties of the cluster are then computed given the stellar population
present at any time. We note that the model output of STARS is quantitatively similar to that from
version 5.1 of Starburst99, which unlike earlier versions does include AGB tracks (Leitherer et al. 1999;
Vazquez & Leitherer 2005). As discussed in detail below, particular predictions of STARS include, for ages
greater than 10 Myr: equivalent widths of WCO2−0 ∼ 12Å and WCO6−3 ∼ 4.5Å, and H-K color of 0.15 mag.
The equivalent quantities predicted by Starburst99 v5.1 are WCO2−0 ∼ 11Å and WCO6−3 ∼ 5Å, and H-K
color of 0.2 mag.
3.1. Isolating the stellar continuum
For small observational apertures a significant fraction of the K-band (and even H-band) continuum
can be associated with non-stellar AGN continuum. The AGN contribution can be estimated from a simple
measurement of the equivalent width of a stellar absorption feature. We use CO 2-0 2.29µm in the K-band or
CO 6-3 1.62µm in the H-band. Although the equivalent widths WCO2−0 and WCO6−3 vary considerably for
individual stars, the integrated values for stellar clusters span only a rather limited range. This was shown
by Oliva et al. (1995) who measured these values for elliptical, spiral, and star-forming (Hii) galaxies. We
have plotted their measurements of these two absorption features in the left-hand panel of Fig. 2, together
with the equivalent widths of giant and supergiant stars from Origlia et al. (1993).
In STARS we use empirically determined equivalent widths from library spectra (Förster Schreiber
2000) to compute the time-dependent equivalent width for an entire cluster of stars. Results for various star
formation histories are shown in the centre and right panels of Fig. 2, for WCO6−3 and WCO2−0 respectively.
Typical values are WCO6−3 = 4.5Å and WCO2−0 = 12Å. The dashed box in the left panel shows that the
– 7 –
locus of 20% deviation from each of these computed values is consistent with observations. That the Hii
galaxies have slightly higher WCO2−0 can be understood because these are selected to have bright emission
lines and hence are strongly biassed towards young stellar ages – often corresponding to the maximum depth
of the stellar features that occurs at 10 Myr due to the late-type supergiant population. It may be this
bias for galaxies selected as ‘starbursts’, and the similarity of the CO depth for starbursts of all other ages,
that led Ivanov et al. (2000) to conclude that there is no evidence for strong starbursts in Seyfert 2 galaxies.
Similarly, an estimate of the dilution can be found from the Nai 2.206µm line. Fig. 7 of Davies et al. (2005)
shows that for nearly all star formation histories the value WNa I remains in the range 2–3Å.
Our conclusion here is that within a reasonable uncertainty of ±20% (see Fig. 2), one can assume that
the intrinsic equivalent width of the absorption – most notably CO – features of any stellar population that
contains late-type stars is independent of the star formation history and age. For a stellar continuum diluted
by additional non-stellar emission, the fraction of stellar light is
fstellar = Wobs/Wint
where Wobs and Wint are the observed and intrinsic equivalent widths of the CO features discussed above.
Thus, we are able to correct the observed continuum magnitude for the contribution associated with the
3.2. Stellar colour and luminosity
Our data cover both the H and K-bands – hence the reason for using both WCO6−3 and WCO2−0. In
order to homogenize the dataset, we need to convert H-band stellar magnitudes to K-band. The STARS
computation in Fig. 3 shows that this conversion is also independent of the star formation history, being
close to H −K = 0.15 mag (no extinction) for all timescales and ages. This result is supported empirically
by photometry of elliptical and spiral galaxies performed by Glass (1984). For ellipticals H −K ∼ 0.2–0.25,
and for spirals H − K ∼ 0.2–0.3. Some of the difference between the data and models could be due to
extinction since H −K = (H −K)0 + (AH −AK); and for AV = 1, AH −AK = 0.08. However, at the level
of precision required here, the 5–10% difference between model and data can be considered negligible.
To convert from absolute magnitude to luminosity we use the relation
MK = −0.33 − 2.5 logLK
where LK is the total luminosity in the 1.9–2.5µm band in units of bolometric solar luminosity (1L⊙ =
3.8 × 1026 W), and as such different from the other frequently used monochromatic definition with units of
the solar K-band luminosity density (2.15 × 1025 Wµm−1). We then use STARS to estimate the bolometric
stellar luminosity Lbol. The relation between Lbol and LK is shown in the right panel of Fig. 3. The
dimensionless ratio Lbol/LK depends on the age and the exponential decay timescale of the star formation.
However, the range spanned is only 20–200 for ages greater than 10 Myr. Thus even if the star formation
history cannot be constrained, a conversion ratio of Lbol/LK ∼ 60 will have an associated uncertainty of
only 0.3 dex. In general we are able to apply constraints on the star formation age, and so our errors will be
accordingly smaller.
– 8 –
3.3. Specific Star Formation Diagnostics
Graphs showing how the diagnostics vary with age and star formation timescale are shown in Fig. 4.
3.3.1. Brγ equivalent width
Once the stellar continuum luminosity is known, an upper limit to the equivalent width of Brγ asssociated
with star formation can be found from the narrow Brγ line flux. In some cases it is possible to estimate what
fraction of the narrow Brγ might be associated with the AGN. This can be done both morphologically, for
example if the line emission is extended along the galaxy’s minor axis; and/or kinematically, for example if
the line shows regions that are broader, perhaps with FWHM a few hundred km s−1, suggestive of outflow.
Even if acounting for the AGN contribution is not possible, one may be able to set interesting upper limits or
even rule out continuous star formation scenarios, and put a constraint on the time since the star formation
was active. This can be seen in the lefthand panel of Fig. 4, which shows for example that for ages less than
109 yrs, continuous star formation scenarios will always have WBrγ > 12 Å.
3.3.2. Supernova rate
We estimate the type ii (core collapse) supernova rate νSN from the radio continuum using the relation
(Condon 1992):
LN (W Hz
−1) = 1.3 × 1023 ν−α(GHz) νSN (yr
where LN is the non-thermal radio continuum luminosity, ν is frequency of the observation and α ∼ 0.8 the
spectral index of the non-thermal continuum. This relation was derived for Galactic supernova remnants;
but a similar one, differing only in having a coefficient of 1.1 × 1023, was derived by Huang et al. (1994) for
M 82. For the 5 GHz non-thermal radio continuum luminosity of Arp 220 (176 mJy, Anantharamaiah et al.
2000) it would lead to a supernova rate of 2.9 yr−1, comfortably within the 1.75–3.5yr−1 range estimated by
Smith et al. (1998) based on the detection of individual luminous radio supernovae. This, therefore, seems
a reasonable relation to apply to starbursts.
We have to be careful, however, to take into account any contribution from the AGN to the radio
continuum. Our premise for the nuclei of Seyfert galaxies is that if the nuclear radio continuum is spatially
resolved (i.e. it has a low brightness temperature) and does not have the morphology of a jet, it is likely
to originate in extended star formation. At the spatial scales of a few parsec or more that we can resolve,
emission from the AGN will be very compact. As a result, we can use the peak surface brightness to
estimate the maximum (unresolved) contribution from an AGN. Wherever possible, we use radio continuum
observations at a comparable resolution to our data to derive the extended emission; and observations at
higher resolution to estimate the AGN contribution. Details of the data used in each case are given in
the relevant sub-sections for each object in Appendix A. In addition, we exclude any emission obviously
associated with jets, for example as in NGC 1068.
To use νSN as a diagnostic, we normalise it with respect to the stellar K-band luminosity. This gives
the ratio 1010νSN/LK , for which STARS output is drawn in Fig. 4.
– 9 –
3.3.3. Mass-to-light ratio
Models indicate that the ratio M/LK of the stellar mass to K-band luminosity should be an excellent
diagnostic since, for ages greater than 10 Myr, it increases monotonically with age as shown in Fig. 4.
However, in practice estimating the stellar mass is not entirely straightforward. In many cases it is only
practicable to derive the dynamical mass. It may be possible to estimate and hence correct for the molecular
gas mass based on millimetre CO maps, but these are scarce at sufficiently high spatial resolution and are
associated with their own CO-to-H2 conversion uncertainties. We also note that it is often not possible to
separate the ‘old’ and ‘young’ stellar populations. The best one can do is estimate the overall mass-to-light
ratio, and argue that this is an upper limit to the true ratio for the young population. While there inevitably
remains uncertainty on the true ratio, the limit is often sufficient to apply useful constraints on the age of
the ‘young’ population.
Our estimates of the dynamical mass are based wherever possible on the stellar kinematics, since the gas
kinematics can be perturbed by warps, shocks, and outflows. We begin by estimating the simple Keplerian
mass assuming that the stars are supported by ordered rotation at velocity Vrot = Vobs/ sin i in a thin plane.
However, the stellar kinematics in all the galaxies exhibit a significant velocity dispersion indicating that a
considerable mass is supported by random rather than ordered motions. Thus the simple Keplerian mass
is very much an underestimate, and any estimate of the actual mass is associated with large uncertainties
– see for example Bender et al. (1992), who derive masses of spheroidal systems. As stated in Section 2,
we assume that the random motions are isotropic. Our relation for estimating the mass enclosed within a
radius R is then
M = (V 2rot + 3σ
2)R/G.
where σ is the observed 1-dimensional velocity dispersion.
We note that when taking rotation into account in estimating the masses of spheroids with various
density profiles, Bender et al. (1992) also use a factor 3 between the V and σ terms in their Appendix B.
Despite the complexities involved, within the unavoidable uncertainties (a factor 2–3), their relation gives
the same mass as that above. Although this uncertainty appears to be quite large, it does not impact the
results and conclusions in this paper since we are concerned primarily with order-of-magnitude estimates
when considering mass surface densities.
4. Properties of Nuclear Star Formation
In the following section we bring to together the individual results (detailed in Appendix A) to form a
global picture. It is possible to do this because all the data have been analysed in a consistent manner, using
the tools described in Section 3 to compare in each object the same diagnostics to the same set of stellar
evolutionary synthesis models.
We note that the discussion that follows is based on results for 8 of the AGN we have observed. As
explained in Appendix A we exclude NGC 2992 because we are not able to put reliable constraints on the
properties of the nuclear star formation. Despite this, there are indications that at higher spatial resolution
one should expect to find a distinct nuclear stellar population as has been seen in other AGN.
– 10 –
Size scale Tracing the stellar features rather than the broad-band continuum, we have in all cases resolved
a stellar population in the nucleus close around the AGN. While this should not be unexpected if the stellar
distribution follows a smooth r1/4 or exponential profile, we have in several cases been able to show that on
scales of < 50 pc there is in fact an excess above what one would expect from these profiles. This suggests
that in general we are probing an inner star forming component.
Fig. 5 shows normalised azimuthally averaged stellar luminosity profiles for the AGN. These have not
been corrected for a possible old underlying population, nor has any deconvolution with the PSF been
performed. Nevertheless, it is still clear that the stellar intensity increases very steeply towards the nucleus.
In 6 of the 8 galaxies shown, the half-width at half-maximum is less than 50 pc. The remaining 2 galaxies
are the most distant in the sample, and the spatial resolution achieved does not permit a size measurement
on these scales. We may conclude that the physical radial size scale of the nuclear star forming regions in
Seyfert galaxies does not typically exceed 50 pc.
Stellar Age For 8 of the AGN studied here, we have been able to use classical star formation diagnostics
based on line and continuum fluxes as well as kinematics to constrain the ages of the inner star forming
regions. The resulting ages should be considered ‘characteristic’, since in many cases there may simultane-
ously be two or more stellar populations that are not co-eval. For example, if a bulge population exists on
these small spatial scales, it was not usually possible to account for the contamination it would introduce.
While this would have little effect on WBrγ , it could impact on M/LK more strongly, increasing the inferred
age. The ages we find lie in the range 10–300Myr, compelling evidence that it is common for there to be
relatively young star clusters close around AGN.
Intriguingly, we also find rather low values of WBrγ : typically WBrγ . 10 Å (see Table 3). This indicates
directly that there is currently little or no on-going star formation. Coupled with the relatively young ages,
we conclude that the star formation episodes are short-lived. One may speculate then that the star formation
is episodic, recurring in short bursts. The scale of the bursts and time interval between them would certainly
have an impact on the fraction of Seyfert nuclei in which observational programmes are able to find evidence
for recent star formation.
Nuclear Stellar Disks The first evidence for nuclear stellar disks came from seeing limited optical
spectroscopy, for which a slight reduction in σ∗ was seen for some spiral galaxies (Emsellem et al. 2001;
Márquez et al. 2003; Shapiro et al. 2003). And there is now a growing number of spiral galaxies – more than
30 – in which the phenomenon has been observed, suggesting that they might occur in 30% or more of disk
galaxies (Emsellem 2006c). The σ∗-drop has been interpreted by Emsellem et al. (2001) as arising from a
young stellar population that is born from a dynamically cold gas component, and which makes a significant
contribution to the total luminosity. This appears to be borne out by N-body and SPH simulations of iso-
lated galaxies (Wozniak et al. 2003), which suggest that although the entire central system will slowly heat
up with time, the σ∗-drop can last for at least several hundred Myr. Indeed, preliminary analysis of optical
integral field data for NGC 3623 suggest that the stellar population responsible for the σ∗-drop cannot be
younger than 1 Gyr (Emsellem 2006b).
Our results provide strong support for the nuclear disk interpretation. In previous work (Davies et al.
2006; Mueller Sánchez et al. 2006), we had argued that in both Circinus and NGC 3227 the inner distributions
were disk-like, albeit thickened. We have now found much more direct evidence for this phenomenon in
NGC 1097 and NGC 1068. In both of these galaxies, we have spatially resolved a σ∗-drop and an excess
– 11 –
stellar continuum over the same size scales. In NGC 1097 this size was ∼0.5′′, corresponding to about 40 pc.
For NGC 1068 these effects were measured out to ∼1′′, equivalent to 70 pc. These are not the scale lengths
of the disks, but simply the maximum radius to which we can detect them. In both cases the mean mass
surface densities are of order Σ =(1-3)×104 M⊙ pc
−2. For an infinitely large thin self gravitating stellar disk,
one can use the expression σ2z = 2πGΣz0 to estimate the scale height. Although this may not be entirely
appropriate, we use it here to obtain a rough approximation to the scale heights, which are 5–20 pc. Thus
while the disks appear to be flattened, they should still be considered thick since the radial extent is only a
few times the scale height.
The impact of nuclear starbursts on the central light profile of galaxies was considered theoretically more
than a decade ago by Mihos & Hernquist (1994). They performed numerical simulations of galaxy mergers
to study the mass and luminosity profiles of the remnants, taking gas into account, and estimating the star
formation rate using a modified Schmidt law. They found that there should be a starburst in the nucleus
which would give rise to an excess stellar continuum above the r1/4 profile of the older stars in the merged
system. Several years ago, compact nuclei were found to be present in a significant fraction of spiral galaxies
(Balcells et al. 2003) as well as Coma cluster dwarf ellipticals (Graham & Guzmán 2003). More recently,
nuclei with a median half-light radius of 4.2 pc have been found in the majority of early-type members of
the Virgo Cluster (Côté et al. 2006); and traced out to ∼1′′, equivalent to ∼ 100 pc, in some of the ‘wet’
merger remnants in that cluster (Kormendy et al. 2007). While the nuclear starbursts in these latter cases
are caused by a merger event, whereas those we are studying arise from secular evolution as gas from the
galaxy disk accretes in the nucleus, there appear to be many parallels in the phenomenology of the resulting
starbursts.
Star Formation Rate It is possible to estimate the bolometric luminosity Lbol∗ of the stars from their K-
band luminosity LK even if one knows nothing about the star formation history. As discussed in Section 3,
this would result in an uncertainty of about a factor 3. The diagnostics in Table 3 and discussions in
Appendix A enable us to apply some constraints to the characteristic age of the star formation. Because
continuous star formation is ruled out by the low WBrγ , we have assumed exponential decay timescales of
τSF = 10–100Myr. We have then used STARS to estimate the average star formation rates. In order to
allow a meaningful comparison between the objects, the rates have been normalised to the same area of
1 kpc2. These are the rates given in Table 3. They are calculated simply as the mass of stars produced
divided by the entire time since the star forming episode began. Because τSF is shorter than the age, the
average includes both active and non-active phases of the starburst. Indeed, for τSF = 10 Myr one would
expect the star formation rate during the active phases to be at least a factor of a few, and perhaps an
order of magnitude, greater. The table shows that on scales of a few hundred parsecs one might expect a
few ×10 M⊙ yr
−1 kpc−2, while on scales of a few tens of parsecs mean rates reaching ∼ 100 M⊙ yr
−1 kpc−2
should not be unexpected; and correspondingly higher – up to an order of magnitude, see Fig. 6 – during
active phases.
An obvious question is why there should be such vigorous star formation in these regions. Star formation
rates of 10–100M⊙ yr
−1 kpc−2 are orders of magnitude above those in normal galaxies and comparable to
starburst galaxies. The answer may lie in the Schmidt law and the mass surface densities we have estimated
in Table 3. Fig. 7 shows these surface densities at the radii over which they were estimated, revealing a trend
towards higher densities on smaller scales and values of a few times 104 M⊙ pc
−2 in the central few tens of
parsecs. The global Schmidt law, as formulated by Kennicutt (1998), states that the star formation rate
depends on the gas surface density as ΣSFR ∝ Σ
gas. If one assumes that 10–30% of the mass in our AGN is
– 12 –
gas, then this relation would predict time-averaged star formation rates in the range 10–100M⊙ yr
−1 kpc−2,
as have been observed. That the high star formation rates may simply be a consequence of the high mass
surface densities is explored futher by Hicks et al. (in prep.).
Stellar Luminosity As a consequence of the high star formation rates, the stellar luminosity per unit area
close around the AGN is very high in these objects. Despite this, because the star formation is occurring
only in very small regions, the absolute luminosities are rather modest. This can be seen in Fig. 8 which
shows the bolometric luminosity of the stars as a fraction of the entire bolometric luminosity of the galaxy.
We have calculated a range for the ratio Lbol∗/LK appropriate for each galaxy based on the ages in Table 3
for different τSF. Because we assume that all the K-band stellar continuum is associated with the young
stars, we have adopted the lower end of each range in an attempt to minimise possible overestimation of
Lbol∗. The resulting values for the ratio used span 30–130, within a factor of 2 of the ‘baseline’ value of 60
given in Section 3. In the central few tens of parsecs, young stars contribute a few percent of the total. But
integrated over size scales of a few hundred parsecs, this fraction can increase to more than 20%. On these
scales, the star formation is energetically significant when compared to the AGN. Such high fractions imply
that on the larger scales the extinction to the young stars must be relatively low. On the other hand, on the
smallest scales where in absolute terms the stellar luminosity is small, there could in general be considerable
extinction even at near infrared wavelengths. In this paper we have not tried to account for extinction since
it is very uncertain. The primary effect of doing so would simply be to increase the stellar luminosity above
the values discussed here.
Fig. 9 shows the stellar bolometric luminosity Lbol∗ integrated as a function of radius. All the curves
follow approximately the same trend, with the luminosity per unit area increasing towards smaller scales and
approaching 1013 L⊙ kpc
−2 in the central few parsecs. This appears to be a robust trend and will not change
significantly even with large uncertainties of a factor of a few. It is remarkable that the luminosity density
of 1013 L⊙ kpc
−2 is that estimated by Thompson et al. (2005) for ULIRGs, which they modelled as optically
thick starburst disks. The main difference between the ULIRG model and the starbursts close around AGN
is the spatial scales on which the starburst occurs.
Based on this model, they argued that ULIRGs are radiating at the Eddington limit for a starburst,
defined as when the radiation pressure on the gas and dust begins to dominate over self-gravity. The limiting
luminosity-to-mass ratio was estimated to be ∼ 500 L⊙/M⊙ by Scoville (2003). He argued that in a star
cluster, once the upper end of the main sequence was populated, the radiation pressure would halt further
accretion on to the star cluster and hence terminate the star formation. Following Thompson et al. (2005),
we apply this definition to the entire disk rather than a single star cluster. For the 1013 L⊙ kpc
−2, this implies
a mass surface density of 2 × 104 M⊙ pc
−2. Comparing these quantities to the AGN we have observed, we
find that on scales of a few tens of parsecs they are an order of magnitude below the Eddington limit. On
the other hand, we have already seen that the low WBrγ indicates that there is little on-going star formation
and hence that the starbursts are short-lived. This is important because short-lived starbursts fade very
quickly. As shown in Fig. 6, for a decay timescale of τSF = 10 Myr, Lbol∗ will have decreased from its peak
value by more than an order of magnitude at an age of 100 Myr. Thus it is plausible – and probably likely
– that while the star formation was active, the stellar luminosity was an order of magnitude higher. In this
case the starbursts would have been at, or close to, their Eddington limit at that time.
The luminosity-to-mass ratio of 500 L⊙/M⊙ associated with the Eddington limit is in fact one that all
young starbursts would exceed if, beginning with nothing, gas was accreted at the same rate that it was
converted into stars. That, however, is not a realistic situation. A more likely scenario, shown in Fig. 10, is
– 13 –
that the gas is already there in the disk. In this case, a starburst with a star-forming timescale of 100 Myr
could never exceed 100 L⊙/M⊙. To reach 500 L⊙/M⊙, the gas would need to be converted into stars on a
timescale . 10 Myr. This timescale is independent of how much gas there is. Thus, for a starburst to reach
its Eddington limit, it must be very efficient, converting a significant fraction of its gas into stars on very
short ∼ 10 Myr timescales. This result is consistent with the prediction of the Schmidt law, which states that
disks with a higher gas surface density will form stars more efficiently. The reason is that the star formation
efficiency is simply SFE = ΣSFR/Σgas ∝ Σ
gas. Thus, from arguments based solely on the Schmidt law and
mass surface density, one reaches the same conclusion that the gas supply would be used rather quickly and
the lifetime of the starburst would be relatively short.
Summarising the results above, a plausible scenario could be as follows. The high gas density leads to
a high star formation rate, producing a starburst that reaches its Eddington limit for a short time. Because
the efficiency is high, the starburst can only be active for a short time and then begins to fade. Inevitably,
one would expect that the starburst is then dormant until the gas supply is replenished by inflow. This
picture appears to be borne out by the observations presented here.
5. Starburst-AGN Connection
In the previous sections we have presented and discussed evidence that in general there appears to have
been moderately recent star formation on small spatial scales around all the AGN we have observed. Fig. 11
shows the first empirical indication of a deeper relationship between the star formation and the AGN. In this
figure we show the luminosity of the AGN, both in absolute units of solar luminosity and also in relative units
of its Eddington luminosity LEdd, against the age of the most recent known nuclear star forming episode.
Since the AGN luminosity is not well known, we have made the conservative assumption that it is equal to
half the bolometric luminosity of the galaxy – as may be the case for NGC 1068 (Pier et al. 1994, but see
also Bland-Hawthorne et al. 1997). To indicate the expected degree of uncertainty in this assertion we have
imposed errorbars of a factor 2 in either direction, equivalent to stating that the AGN luminosity in these
specific objects is likely to be in the range 25–100% of the total luminosity of the galaxy. The Eddington
luminosity is calculated directly from the black hole mass, for which estimates exist for these galaxies from
reverberation mapping, the MBH − σ∗ relation, maser kinematics, etc. These are listed in Table 2. For the
age of the star formation, we have plotted the time since the most recent known episode of star formation
began, as given in Table 3. For galaxies where a range of ages is given, we have adopted these to indicate
the uncertainty; the mean of these, ∼ ±30%, has been used to estimate the uncertainty in the age for the
rest of the galaxies. We note that these errorbars reflect uncertainties in characterising the age of the star
formation from the available diagnostics and also in the star formation timescale τSF. However, there are still
many implicit assumptions in this process, and we therefore caution that the actual errors in our estimation
of the starburst ages may be larger than that shown.
Conceding this, we do not wish to over-interpret the figure. Keeping the uncertainties in mind, Fig. 11
shows the remarkable result that AGN which are radiating at lower efficiency . 0.1 L/LEdd are associated
with younger . 50–100Myr starbursts; while those which are more efficient & 0.1 L/LEdd have older &
50–100Myr starbursts. If one were to add to this figure the Galactic Centre – which is known to have
an extremely low luminosity (L/LEdd < 10
−5; Ozernoy & Genzel 1996; Baganoff et al. 2003) and to have
experienced a starburst 6 ± 2 Myr ago (Paumard et al. 2006) – it would be consistent with the categories
above. The inference is that either there is a delay between the onset of starburst activity and the onset of
AGN activity, or star formation is quenched once the black hole has become active.
– 14 –
In Section 4 we argued that the starbursts are to some extent self-quenching: that very high star forma-
tion efficiencies are not sustainable over long periods. In addition, an intense starburst will provide significant
heat input to the gas, which is perhaps partially responsible for the typically high gas velocity dispersions
in these regions (Hicks et al., in prep.). This itself could help suppress further star formation. Heating by
the AGN could also contribute to this process, and has been proposed as the reason why the molecular
torus is geometrically thick (Pier & Krolik 1992; Krolik 2007). It is also used to modulate star formation (at
least on global scales) in semi-analytic models of galaxy evolution (Granato et al. 2004; Springel et al. 2005).
While this is certainly plausible, it does not explain either why the star formation in some galaxies with a
lower luminosity AGN has already ceased, nor why none of the AGN associated with younger starbursts are
accreting efficiently.
Instead we argue for the former case above, that efficient fuelling of a black hole is associated with a
starburst that is at least 50–100Myr old. It may be because of such a delay between AGN and starburst
activity that recent star formation is often hard to detect close to AGN: the starburst has passed its most
luminous (very young) age, and is in decline while the AGN is in its most active phase (see Fig. 6). This
does not necessarily imply that the a priori presence of a starburst is required before an AGN can accrete
gas – although it seems inevitable that one will occur as gas accumulates in the nucleus. Nor does it imply
that all starbursts will result in fuelling a black hole; indeed it is clear that there are many starbursts not
associated with AGN. As we argue below, the crucial aspect may be the stellar ejecta associated with the
starburst; and in particular, not just the mass loss rate, but the speed with which the mass is ejected.
Winds from OB stars In the Galactic Center, Ozernoy & Genzel (1996) proposed that it is the recent
starburst there that is limiting the luminosity of the black hole. In this scenario, mechanical winds from
young stars – both the outflow and the angular momentum of the gas (which is a consequence of the angular
momentum of the stars themselves) – hinder further inflow. The authors argued that almost none of the
gas flowing into the central parsec reached the black hole because of outflowing winds from IRS 16 and He i
stars in that region. Detailed modelling of the Galactic Center region as a 2-phase medium was recently
performed by Cuadra et al. (2006). They included both the fast young stellar winds with velocities of
700 km s−1 (Ozernoy et al. 1997) and the slower winds of ∼ 200 km s−1 (Paumard et al. 2001); and also took
into account the orbital angular momentum of the stars (Paumard et al. 2001; Genzel et al. 2003), which
had a strong influence on reducing the accretion rate. They found that the average accretion rate onto
the black hole was only ∼ 3 × 10−6 M⊙ yr
−1, although an intermittent cold flow superimposed considerable
variability onto this. In contrast, the hypothetical luminosity Ozernoy & Genzel (1996) estimate that Sgr A∗
would have if it could accrete all the inflowing gas, would be 5 × 1043 erg s−1, typical of Seyfert galaxies.
In principle this process could be operating in other galaxy nuclei where there has been a starburst which
extends to less than 1 pc from the central black hole. However, it cannot explain the timescale of the delay we
have observed, which is an order of magnitude greater than the main sequence lifetime of OB and Wolf-Rayet
stars.
Winds from AGB stars Stars of a few (1–8 M⊙) solar masses will evolve on to the asymptotic giant
branch (AGB) at the end of their main sequence lifetimes. The timescale for stars at the upper end of this
range to reach this phase is ∼ 50 Myr, comparable to the delay apparent in Fig. 9. Since AGB stars are known
to have high mass-loss rates, of order 10−7–10−4 M⊙ yr
−1 at velocities of 10–30km s−1 (Winters et al. 2003),
they may be prime candidates for explaining the delay between starburst and AGN activity. To quantify
this, we consider how much of the mass in the wind could be accreted by the central supermassive black
– 15 –
hole.
The Bondi parameterisation of the accretion rate onto a point particle for a uniform spherically sym-
metric geometry is given by (Bondi 1952)
2πG2 M2 ρ
(V 2 + c2s)
where M is the mass of the point particle moving through a gas cloud, V is the velocity of the particle with
respect to the cloud, ρ is the density of the cloud far from the point particle, and cs is the sound speed.
This approximation is still used to quantify accretion on to supermassive black holes in models of galaxy
evolution (Springel et al. 2005), even though it may be significantly inaccurate for realistic (e.g. turbulent)
media (Krumholz et al. 2006). Here, it is sufficient to provide an indication of the role that stellar winds
may play in accretion onto a central black hole. The density of the stellar wind at a distance R from the
parent star is given by
ρwind =
Ṁwind
4R2 Vwind
In our case, R is the distance from the star to the black hole. One would therefore expect that the accretion
rate on to the black hole could be written as (see also Melia 1992)
ṀBH ∼
G2 M2BH Ṁwind
+ c2s)
3/2 Vwind R2
This equation shows that ṀBH ∝ V
. We have implicitly assumed that Vwind is greater than the orbital
velocity Vorb of the star from which it originates. This is not the case for AGB winds, and so one reaches
the limiting case of ṀBH ∝ V
, where for the galaxies we have observed Vorb ∼ 50–100km s
−1. This is still
at least an order of magnitude less than the winds from OB and Wolf-Rayet stars. Thus, even though the
mass loss rates from individual OB and Wolf-Rayet stars are similar to those of AGB stars, the AGB winds
will fuel a black hole much more efficiently. However, for slow stellar winds that originate close to a 107 M⊙
black hole, the equation breaks down because the conditions of uniformity and spherically symmetry are
strongly violated. Indeed, the apparent accretion rate exceeds the outflow rate – implying that essentially
the entire wind can be accreted. For AGB wind velocities of 10–30 km s−1, the maximum radius at which
the entire wind from a star in Keplerian orbit around a 107 M⊙ black hole will not exceed the escape velocity
from that orbit (i.e. Vwind + Vorb < Vesc) is around 10–70pc. We adopt the middle of this range, 40 pc, as
the characteristic radius within which it is likely that a significant fraction, and perhaps most, of the AGB
winds are accreted onto the black hole. Fig. 9 indicates that the stellar luminosity within this radius is
∼ 2 × 109 L⊙. It is this luminosity that has been used to scale the STARS model (for τSF = 10 Myr and an
age of 100 Myr) in Fig. 6, and so one can also simply read off the mass loss from the figure. The mass loss
rate for such winds peaks at about 0.1 M⊙ yr
−1 and then tails off proportionally to the K-band luminosity,
leading to a cumulative mass lost of 2× 107 M⊙ after 1 Gyr (although most of the loss occurs actually occurs
within half of this timespan). This mass loss rate is sufficient to power a Seyfert nucleus for a short time. A
typical Seyfert with MBH ∼ 10
7 M⊙ requires 0.02 M⊙ yr
−1 to radiate at the Eddington limit. Even for the
short bursts we have modelled, Fig. 6 shows that this can be supplied by AGB winds for starburst ages in
the range 50–200Myrs.
We note that taking an AGB star luminosity of 104 L⊙ (which is at the high end of the likely average,
Nikolaev & Weinberg 1997) we then find that there are ∼ 2 × 105 AGB stars close enough to the black hole
to contribute to accretion. In order to provide at least 0.02 M⊙ yr
−1, the typical mass loss rate per star must
– 16 –
exceed 10−7 M⊙ yr
−1, which is the lower limit of the range measured for Galactic AGB stars given above.
Thus the mass losses and rates estimated here appear to be plausible.
The low speed of these winds means they will not create much turbulence. We quantify this by consid-
ering their total mechanical energy 1
mv2 integrated over the same timespan, which is ∼ 1045 J. These two
quantities – gas mass ejected and mechanical energy – are compared to those for supernovae below.
Supernovae Type ii supernovae are the stellar outflows most able to create turbulence in the interstellar
medium, since they typically eject masses of ∼ 5 M⊙ at velocities of ∼ 5000 km s
−1 (Chevalier 1977). Each
supernova therefore represents a considerable injection of mechanical momentum and energy into the local
environment. A large number of compact supernova remnants are known, for example in M 82 and Arp 220,
and are believed to have expanded into dense regions with nH ∼ 10
3–104 cm−3 (Chevalier & Fransson 2001).
These authors argue that such remnants become radiative when they reach sizes of ∼ 1 pc, at which point
the predicted expansion velocity will have slowed to ∼ 500 km s−1. By this time, the shock front will have
driven across ∼ 1000 M⊙ of gas. When integrated over the age of the starburst, even for low supernova
rates – e.g. the current rate within 30 pc of the nucleus of NGC 3227 is ∼ 0.01 yr−1 (Davies et al. 2006)
– this represents a substantial mass of gas that has been affected by supernova remnants. The STARS
model we have constructed in Fig. 6 indicates that typically one could expect ∼ 106 supernovae to occur as
a result of one of the short-lived starbursts; and that most of these will occur around 10–50Myr after the
beginning of the starburst. For a decay timescale of the star formation rate that is longer than τSF = 10 Myr,
this timespan will increase. Hence, supernovae may also play a role in causing the observed delay between
starburst and AGN activity.
STARS calculates the mass loss and mass loss rates using a very simple scheme, assuming that a star
ejects all of its lost mass at the end of its life on a stellar track. Thus, it does not calculate the mass lost
from supernovae explicitly, rather the combined mass lost from OB winds and supernovae which is much
higher. We therefore adopt the ∼ 5 M⊙ per supernova given above, which yields a total ejected mass of
∼ 8 × 106 M⊙. This is about 40% of that released by AGB winds. However, since this gas is ejected at high
speed and ṀBH ∝ V
, the efficiency with which it can be accreted onto the black hole is extremely low.
This can also be seen in the total mechanical energy of ∼ 1050 J, which is several orders of magnitude greater
than for AGB winds. In fact the total mechanical energy exceeds the binding energy of the nuclear region,
which is of order 1048 J (assuming 108 M⊙ within 40 pc). As a result, it is highly likely that supernova cause
some fraction of the gas to be permanently expelled. Indeed, superwinds driven by starbursts are well known
in many galaxies. This is not important as long as sufficient gas either remains to fuel the AGN, or more is
produced by stellar winds – which, as we have argued above, appears to be the case for AGB stars.
6. Conclusions
We have obtained near infrared spectra of 9 nearby active galactic nuclei using adaptive optics to achive
high spatial resolution (in several cases better than 10 pc). For 7 of these, integral field spectroscopy with
SINFONI allows us to reconstruct the full 2-dimensional distributions and kinematics of the stars and gas.
Although the individual AGN are very varied, we have analysed them in a consistent fashion to derive: the
stellar K-band luminosity, the dynamical mass, and the equivalent width of the Brγ line. We have combined
these with radio continuum data from the literature, which has been used to estimate the supernova rate.
We have used these diagnostics to constrain STARS evolutionary synthesis models and hence characterize
– 17 –
the star formation timescales and ages of the starbursts close around AGN. Our main conclusions can be
summarised as follows:
• The stellar light profiles show a bright nuclear component with a half-width at half-maximum of
less than 50 pc. In a number of cases these nuclear components clearly stand out above an inward
extrapolation of the profile measured on larger scales. In addition, there are 2 cases which show
kinematical evidence for a distinct stellar component, indicating that the nuclear stellar populations
most probably exist in thick nuclear disks. The mean mass surface densities of these disks exceeds
104 M⊙ pc
• There is abundant evidence for recent star formation in the last 10–300Myr. But the starbursts are no
longer active, implying that the star formation timescale is short, of order a few tens of Myr. While the
starbursts were active, the star formation rates would have been much higher than the current rates,
reaching as high as 1000 M⊙ kpc
−2 in the central few tens of parsecs (comparable to ULIRGs, but on
smaller spatial scales). These starbursts would have been Eddington limited. Due to the very high star
forming efficiency, the starbursts would have also exhausted their fuel supply on a short timescale and
hence have been short-lived. It therefore seems likely that nuclear starbursts are episodic in nature.
• There appears to be a delay of 50–100Myr (and in some cases perhaps more) between the onset of star
formation and the onset of AGN activity. We have interpreted this as indicating that the starburst
has a significant impact on fuelling the central black hole, and have considered whether outflows from
stars might be responsible. While supernovae and winds from OB stars eject a large mass of gas, the
high velocity of this gas means that its accretion efficiency is extremely low. On the other hand, winds
from AGB stars ultimately dominate the total mass ejected in a starburst; and the very slow velocities
of these winds mean they can be accreted onto the black hole very efficiently.
The authors thank all those who assisted in the observations, and also the referee for a thorough review
of the paper. This work was started at the Kavli Institute for Theoretical Physics at Santa Barabara and
as a result was supported in part by the National Science Foundation under Grant No. PHY05-51164. RD
aknowledges the interesting and useful discussions he had there with Eliot Quartaert, Norm Murray, Julian
Krolik and Todd Thompson.
Facilities: Keck:II (NIRSPAO, NIRC2), VLT:Yepun (NACO, SINFONI).
A. Individual Objects
This appendix contains specific details on the individual objects. We summarize our published results
from near infrared adaptive optics spectroscopy of individual objects, and present a brief analysis of the
new data for several other objects. The aim of re-assessing the data for Mkn 231 that has already been
published is to ensure that it is analysed using STARS in a manner that is consistent with the new data. For
NGC 7469, we make a significant update of the analysis using new data from integral field spectroscopy. In
general, for objects with new data, we provide only the part of the analysis relevant to understanding star
formation around the AGN. Our intention is that a complete analysis for each object will be presented in
future publications.
– 18 –
Our analyses are restricted to the nuclear region. Since there is no strict universal definition of what
comprises the ‘nuclear region’, we explicitly state in Table 3 the size of the region we study in each galaxy.
The table also presents a summary of the primary diagnostics. The way in which these have been derived,
and their likely uncertainties, has been discussed in some detail already in Section 3. As such, the description
of these methods is not repeated, and in this Section we discuss only issues that require special attention.
A.1. Summary of Star Forming Properties of Galaxies already Studied
A.1.1. Mkn 231
A detailed analysis of the star formation in the nucleus of Mkn 231 at a resolution of about 0.18′′ (150 pc)
was given in Davies et al. (2004b). Here we summarize only the main points; no new data is presented, but
the analysis is updated using STARS to make it consistent with the other objects studied in this paper.
The presence of stellar absorption features across the nucleus demonstrates the existence of a significant
population of stars. The radial distribution and kinematics indicate they lie, like the gas (Downes & Solomon
1998), in a nearly face-on disk. Davies et al. (2004b) found that the dynamical mass imposed a strong
constraint on the range of acceptable starburst models, yielding an upper limit to the age of the stars of
around 120 Myr. Re-assessing the mass-to-light ratio using STARS models suggests that for the increased
mass required by a more face-on orientation (i = 10◦) an upper age of 250 Myr is also possible, depending
on the star formation timescale. However, either a small change of only a few degrees to the inclination (to
i = 15◦), or a relatively short star formation timescale of 10 Myr would reduce the limit to the ∼100 Myr
previously estimated. This is more consistent with the extremely high supernova rate.
The stellar luminosity, found from the dilution of the CO absorption (Davies et al. 2004b), indicates
that stars within 1′′ (800 pc) of the nucleus contribute 25–40% of the bolometric luminosity of the galaxy.
Similarly, within 200 pc, stars comprise 10–15% of Lbol. The age, star formation rate, and size scale (disk
scale length of 0.18–0.2′′) are all consistent with high resolution radio continuum imaging (Carilli et al. 1998).
A.1.2. Circinus
Star formation in the central 16 pc of Circinus was addressed by Mueller Sánchez et al. (2006). The
diagnostics given in Table 3 are taken from this reference. We used the depth of the CO 2-0 bandhead to
estimate the stellar luminosity, combined with the narrow Brγ flux (which we argued originated in star
forming regions rather than the AGN narrow line region) and the radio continuum, to constrain starburst
models. The conclusion was that the starburst was less than 80 Myr old and was already decaying. On
these scales it contributes 1.4% of Lbol, or more if extinction is considered. A similar nuclear star formation
intensity was estimated by Maiolino et al. (1998), who were also able to study Circinus on larger scales.
They found that the luminosity of young stars within 200 pc of the AGN was of order 1010L⊙, and hence
comparable to the AGN.
– 19 –
A.1.3. NGC 3227
An analysis similar to that for Circinus was performed on NGC 3227 by Davies et al. (2006), and the
diagnostics given in Table 3 are taken from this reference. In this case we were able to make estimates of
and correct for contributions of: (1) the narrow line region to Brγ, because there were clear regions along
the minor axis that had higher dispersion; (2) the AGN to the radio continuum, by estimating the maximum
contribution from an unresolved source; and (3) the bulge stars to the stellar luminosity, by extrapolating
the radial profile of the bulge to the inner regions. The STARS models yielded the result that in the nucleus,
star formation began approximately 40 Myr ago and must have already ceased. At the resolution of 0.085′′,
the most compact component of stellar continuum had a measured FWHM of 0.17′′, suggesting an intrinsic
size scale of ∼ 12 pc. Young stars within 30 pc of the AGN (i.e. more than just the most compact region)
have a luminosity of ∼ 3 × 109L⊙, which is ∼ 20% of the entire galaxy.
A.2. Star Forming Properties of Galaxies with New Data
A.2.1. NGC 7469
Star formation on large scales in NGC 7469 has been studied by Genzel et al. (1995). They found that
within 800 pc of the nucleus, a region that includes the circumnuclear ring, the luminosity from young stars
was ∼ 3 × 1011L⊙, about 70% of the galaxy’s bolometric luminosity. This situation is similar to that in
Mkn 231. On smaller scales, the nuclear star formation in NGC 7469 was directly resolved by Davies et al.
(2004a) on a size scale of 0.15–0.20′′ (50–65 pc) FWHM. An analysis of the longslit data, similar to that
for Mkn 231, was made – making use of stellar absorption features, kinematics, and starburst models. We
estimated that the age of this region was no more than 60 Myr under the assumption that the fraction of
stellar light in the K-band in the central 0.2′′ was 20–30%. Our new integral field SINFONI observations of
NGC 7469 at a spatial resolution of 0.15′′ (measured from both the broad Brγ and the non-stellar continuum
profiles, see Section 2) are used here to make a more accurate estimate of the nuclear K-band luminosity.
They enable us to provide a short update to the detailed analysis in Davies et al. (2004a).
The SINFONI data show that the equivalent width of the 2.3µm CO 2-0 is WCO2−0 = 1.8Å in a 0.8
aperture and 0.9Å in a 0.2′′ aperture. The corresponding K-band magnitudes are K = 10.4 and K = 11.8
respectively. If one takes the intrinsic equivalent width of the 2.3µm CO 2-0 bandhead to be 12Å (see
Section 3), one arrives at a more modest value of 8% for the stellar fraction of K-band continuum in the
0.2′′ aperture. The stellar K-band luminosity in this region is then 6 × 107 L⊙. Comparing this to the
dynamical mass in Davies et al. (2004a) yields a mass-to-light ratio of M/LK ∼ 0.6 M⊙/L⊙. Previously,
extrapolation from a 37 mas slit to a filled aperture had led to an underestimation of the total magnitude
but an overestimation of the stellar contribution. Fortuitously, these uncertainties had compensated each
other. The same analysis for the 0.8′′ aperture yields a K-band stellar luminosity of 3 × 108 L⊙ and hence
M/LK ∼ 1.6 M⊙/L⊙.
The K-band datacube yields estimates of the upper limit to WBrγ of 17Å and 11Å in 0.2
′′ and 0.8′′
apertures respectively. This has been corrected for dilution of the stellar continuum (as described in Section 3)
but not for a possible contribution to the narrow Brγ from the AGN. Hence the actual WBrγ corresponding
to only the stellar line and continuum emission will be less than these values – indicating that the star
formation is unlikely still to be on-going.
We estimate the age of the star formation using the STARS models in Fig. 4. Within the 0.2′′ aperture
– 20 –
this gives 100 Myr, comparable to our original estimate. Such a young age is supported by radio continuum
measurements. With a 0.2′′ beam, Colina et al. (2001) reported that the unresolved core flux in NGC 7469
was 12 mJy at 8.4 GHz. With much higher spatial resolution of 0.03′′, Sadler et al. (1995) reported an upper
limit to the unresolved 8.4 GHz continuum of 7 mJy. We assume that the difference of 5 mJy is due to
emission extended on scales of 10–60 pc which is resolved out of one beam but not the other. As discussed
in Section 3, star formation is a likely candidate for such emission. In this case, we would estimate the
supernova rate to be ∼ 0.1 yr−1 and the ratio 1010 νSN/LK ∼ 3. This is likely to be a lower limit since there
was only an upper limit on the core radio flux density. For a ratio of this order, even allowing for some
uncertainty, Fig 4 implies an age consistent with no more than 100 Myr.
Within the 0.8′′ aperture, which we adopt in Table 3, continuous star formation is inconsistent with
WBrγ . For a star formation timescale of τSF = 100 Myr, the mass-to-light ratio implies an age of 190 Myr,
just consistent with the measured value of WBrγ = 11Å. If some of the narrow Brγ is associated with the
AGN rather than star formation, then a shorter star formation timescale is required. For τSF = 10 Myr, the
ratio M/LK yields an age of 110 Myr.
A.2.2. IRAS 05189-2524
Fig. 12 shows the H-band spectrum integrated across two segments of the NACO slit, located on either
side of the nucleus. It shows that even away from the nucleus, the depth of the stellar absorption features
is only a few percent. We have therefore decomposed the data into the stellar and non-stellar parts using
both the stellar absorption features and the spectral slope of the continuum. The latter method has been
shown to work for well sampled data by Davies et al. (2004a). The rationale is that the hot dust associated
with the AGN will be much redder than the stellar continuum. An AGN component is also expected to be
unresolved for a galaxy at the distance (170 Mpc) of IRAS 05189-2524. The spectral slope was determined
by fitting a linear function to the spectrum at each spatial position along the slit. It is plotted as a function
of position in Fig. 13, showing a single narrow peak. A Gaussian fit to this yields a spatial resolution of 0.12′′
(100 pc) FWHM. The stellar continuum, also shown in Fig. 13, has been determined by summing the four
most prominent absorption features: CO 4-1, Si I, CO 5-2, CO 6-3. While a Gaussian is not an optimal fit to
this profile, it does yield an aproximate size scale, which we find to be 0.27′′ FWHM. Quadrature correction
with the spatial resolution yields an intrinsic size of 0.25′′ (200 pc). As a cross-check, in the figure we have
compared the sum of these two components to the full continuum profile. The good match indicates that
the decomposition appears to be reasonable.
Remarkably, the 200 pc size of the nuclear stellar light is very similar to that of the 8.44 GHz radio con-
tinuum map of Condon et al. (1991). With a beam size of 0.50′′×0.25′′, they resolved the nuclear component
to have an intrinsic size of 0.20′′×0.17′′. In constrast to radio sources which are powered by AGN and have
brightness temperatures Tb ≫ 10
5 K, the emission here is resolved and has a low brightness temperature
of ∼ 4000 K. This implies a star forming origin. Using their scaling relations further suggests that the flux
density corresponds to a supernova rate of ∼ 1 yr−1.
As described in Section 3, we have estimated the stellar luminosity by comparing the H-band spectrum
to a template star to correct for dilution. We used HR 8465 a K 1.5 I star for which the equivalent width
of CO 6-3 is 4.2Å, within the 4–5Å range predicted by STARS in Fig. 2. By extrapolating from the spatial
profiles along the slit we have estimated the integrated equivalent width within a 1.1′′ aperture, for which
Scoville et al. (2000) gave an H-band magnitude of 11.83. Using all four features above we find for the
– 21 –
template W = 14.4Å and for IRAS 05189-2524 WCO6−3 = 6.7Å. This implies that in the central 1.1
approximately 45% of the H-band continuum originates in stars. Using the colour conversion H −K = 0.15
from Fig. 3 (see Section 3) we find a K-band magnitude for the stars of 12.55 mag and hence a K-band stellar
luminosity of 2×109 L⊙. Putting these results together we derive a ratio of supernova rate to K-band stellar
luminosity of νSN[yr
−1]/LK [10
10L⊙] ∼ 5. Applying corrections for extinction and an AGN contribution
would tend to decrease this ratio.
As a second diagnostic we use WBrγ . We estimate the dilution of the K-band continuum via two
methods. Firstly, we measure WNaI = 0.3Å, indicating a stellar fraction of 0.10–0.15. A consistency check is
provided by the H-band dilution, which we extrapolate to the K-band using blackbody functions for the stars
and dust assuming characteristic temperatures of 5000 K and 1000 K respectively. This method suggests the
K-band stellar fraction is around ∼ 0.14. Hence correcting the directly measured equivalent width of the
narrow Brγ for the non-stellar continuum yields WBrγ = 4–5Å.
Since IRAS 05189-2524 is close to face-on (Scoville et al. 2000), it is not straightforward to make a
reliable estimate of the dynamical mass. Nevertheless, requiring νSN/LK to be high while WBrγ is low
already puts significant constraints on the star formation history. Thus, although the star formation has
probably ended, the age is unlikely to be greater than 100 Myr, and could be as low as 50 Myr where νSN/LK
peaks. For such ages the ratio Lbol/LK is in the range 100–150. Hence for the young stars within 0.55
(450 pc) of the nucleus we find Lbol ∼ (2–3) × 10
11 L⊙, about 20% of Lbol for the galaxy.
A.2.3. NGC 2992
The spatial resolution of the K-band data for NGC 2992 has been estimated from both the broad Brγ
and the non-stellar continuum (see Section 2 and 3). The two methods yield symmetric PSFs, with FWHMs
of 0.32′′ and 0.29′′ respectively, corresponding to 50 pc.
The CO 2-0 equivalent width of ∼ 3Å implies a stellar fraction of ∼ 0.25 within a radius of 0.4′′, and
hence a stellar luminosity of LK = 3.5 × 10
7 L⊙.
Unlike IRAS 05189-2524, the radio continuum in NGC 2992 is quite complex. Much of the extended
emission on scales of a few arcsec appears to originate from a superbubble, driven either by the AGN or
by a nuclear starburst. On the other hand, most of the nuclear emission seems to be unresolved. With a
beam size of 0.34′′×0.49′′, Wehrle & Morris (1988) measured the unresolved flux to be 7 mJy at 5 GHz. At
a resolution better than 0.1′′, Sadler et al. (1995) reported a 2.3 GHz flux of 6 mJy. Based on this as well as
non-detections at 1.7 GHz and 8.4 GHz, they estimated the core flux at 5 GHz to be <6 mJy. Taking a flat
spectral index, as indicated by archival data (Chapman et al. 2000), one might expect the 5 GHz core flux
to be not much less than 6 mJy, leaving room for only ∼ 1 mJy in extended emission in the central 0.5′′. If
we assume this difference can be attributed to star formation, it implies a supernova rate of ∼ 0.003 yr−1
and hence 1010 νSN/LK ∼ 1. Fig. 4 shows that a ratio of this order is what one might expect for ages up to
200 Myr. However, given the uncertainty it does not impose a significant constraint.
It is also difficult to quantify what fraction of the narrow Brγ is associated with star formation. This is
made clear in Fig. 14 which shows that the morphology of the line (centre left panel) does not follow that of
the stars (far left). In addition, particularly the south-west side is associated with velocities that are bluer
than the surrounding emission, indicative of motion towards us. The western edge also exhibits high velocity
dispersion. Taken together, these suggest that we may be seeing outflow from the apex of an ionisation cone
– 22 –
with a relatively large opening angle. This interpretation would tend to support the hypothesis that the
radio bubble has been driven by the AGN.
The stellar continuum appears to trace an inclined disk, the north west side of which is more obscured
(Fig. 14). However, the velocity dispersion is high, exceeding 150 km s−1 across the whole field (Fig. 15). This
is similar to the 160 km s−1 reported by Nelson & Whittle (1995) from optical spectroscopy, and suggests
that we are seeing bulge stars. To analyse the radial luminosity profile we have fitted it with both an r1/4
law and exponential profile. The fits in Fig. 16 were optimised at radii r > 0.5′′ and then extrapolated
inwards, convolved with the PSF. Whether one could claim that there is excess continuum in the nucleus
depends on the profile fitted. The r1/4 law provides a stronger constraint since it is more cuspy, and suggests
there is no excess. Although this evidence is inconclusive, Fig. 15 suggests that there is some kinematic
evidence favouring the existence of a distinct nuclear stellar population. This comes in the form of a small
unresolved drop in dispersion at the centre, similar to those in NGC 1097 and NGC 1068. While the evidence
in NGC 2992 is not compelling, the dispersion is consistent with there being an equivalent – but fainter –
nuclear disk on a scale of less than our resolution of 50 pc. In general it seems that the K-band light we are
seeing is dominated by the bulge, and we are therefore unable to probe in detail the inner region where it
seems that more recent star formation has probably occurred.
Thus, although the available data suggest there has likely been recent star formation in the nucleus of
NGC 2992, the only strong constraint we can apply is that continuous star formation in the central arcsec
over the last Gyr can be ruled out since it would require WBrγ > 10–15Å. We therefore omit NGC 2992 from
the discussion and analysis in Sections 4 and 5.
A.2.4. NGC 1097
In NGC 1097, the first evidence for recent star formation near the nucleus was in the form of a reduction
in the stellar velocity dispersion. Emsellem et al. (2001) proposed this could be explained by the presence of
a dynamically cold nuclear disk that had recently formed stars. Direct observations of a spiral structure in
the central few arcsec, from K-band imaging (Prieto et al. 2005) and [N ii] streaming motions (Fathi et al.
2006), have since confirmed this idea. However, some issues remain open, such as why there are three spiral
arms rather than the usual two, and why gas along one of them appears to be outflowing.
Our data, at a resolution of 0.25′′ measured from the H-band non-stellar continuum, also reveal the
same spiral structure. Indeed, we find that it is traced by the morphology of the CO bandhead absorption as
well as by the 2.12µm H2 1-0 S(1) line. Interestingly, 1-0 S(1) emission is stronger where the stellar features
are weaker. This suggests that obscuration by gas and dust plays an important role. Fig. 18 shows that an
r1/4 law, typical of stellar bulges, with effective radius Reff = 0.5
′′ is a good fit to the stellar radial profile
at 0.5′′ < r < 1.8′′. It therefore seems reasonable to argue that at these radii it is only the gas that lies in a
disk. In this picture the spiral structure in the stellar continuum arises solely due to extinction of the stars
behind the disk. Extrapolating this fit, convolved with the PSF, to the nucleus indicates that at r < 0.5′′
there is at least 25% excess stellar continuum. There could be much more, given that it coincides with a
change in the dominant kinematics.
For NGC 1097 we parameterized the kinematics of the gas and stars quantitatively using kinemetry.
Based on the uniformity of the velocity field, we made the simplifying assumption that across the central 4′′
the gas lies in a single plane whose centre is coincident with the peak of the non-stellar emission. We were then
able to derive the position angle and inclination of the disk (see Section 2). The 2D kinematics of the stars
– 23 –
is traced via the CO2-0 absorption bandhead, and that of the gas through the 1-0S(1) emission line. These
independently yielded similar parameters: both gave a position angle of −49◦ and their inclinations were 43◦
and 32◦ respectively. These are fully consistent with values found by other authors (Storchi-Bergmann et al.
2003; Fathi et al. 2006). The resulting rotation curves and velocity dispersions are shown in Fig. 19. The
residuals, which can be seen in the velocity field of the gas but not the stars, and their relation to the spiral
structure described above will be discussed elsewhere (Davies et al. in prep). The important result here is
that at our spatial resolution, we find that the central stellar dispersion is σ∗ = 100 km s
−1, less than the
surrounding 150 km s−1 and also less than that in the seeing limited spectra of Emsellem et al. (2001). In
the same region we find that the rotation velocity of the gas starts to decrease rapidly, and its dispersion
increases from σgas ∼ 40 km s
−1 to ∼ 80 km s−1.
Fig. 19 also shows that while the kinematics of the stars and gas are rather different at large (> 0.5′′)
radii, they are remarkably similar at radii < 0.5′′. This certainly provides a strong indication that in the
nuclear region the stars and gas are coupled, most likely in a (perhaps thick) disk; and that the stars in this
disk, which are bright and hence presumably young, give rise to the excess stellar continuum observed.
Evidence for a recent starburst has been found by Storchi-Bergmann et al. (2005) through optical and
UV spectra. They argued that a number of features they observed could only arise from an 106 M⊙ instanta-
neous starburst, which occurred a few Myr ago and is reddened by AV = 3 mag of extinction. Using STARS
we have modeled this starburst as a 106 M⊙ burst beginning 8 Myr ago with an exponential decay timescale
of 1 Myr. The age we have used is a little older to keep the Brγ equivalent width low; and at this age, the
model predicts WBrγ = 4Å. As Fig. 20 shows, the observed Brγ is weak, although perhaps slightly resolved.
Corrected for the non-stellar continuum, we measure only WBrγ ∼ 1Å. However, the bulge population may
account for a significant fraction of the K-band stellar continuum. Correcting also for this could increase
WBrγ to 2–5Å, consistent with that of the model – assuming that the Brγ is associated with the starburst
rather than the AGN. To within a factor of a few, the scale of the model starburst is also consistent with
that measured: In the central 0.5′′ we measure a Brγ flux of 2 × 10−19 W m−2, compared to that predicted
by the model of 5× 10−19 W m−2. Given the uncertainties – factors of a few – both in the parameters of the
starburst model and also in the corrections we have applied to the data, we consider this a good agreement.
We cannot constrain the starburst further due to its compactness. Storchi-Bergmann et al. (2005) found
that it was occurring in the central 0.2′′, whereas our resolution is only 0.25′′. The Brγ emission is confined
to the central 0.4–0.5′′, although its size is hard to measure due to its weakness with respect to the stellar
absorption features. In this region the K-band stellar luminosity is 4.5× 106 L⊙. To estimate the dynamical
mass we use the mean kinematics of the stars and gas, i.e. Vrot = 40 km s
−1 (corrected for inclination) and
σ = 90 km s−1 (this is the central value, which is least biassed by bulge stars), yielding 1.4× 108 M⊙. This is
actually dominated by the black hole, which has a mass of (1.2±2)×108 M⊙ (Lewis & Eracleous 2006). The
difference between these implies a mass of gas and stars of ∼ 2× 107 M⊙, although with a large uncertainty.
The associated mass-to-light ratio is M/LK ∼ 4. On its own, this implies that over the relatively large area
that it encompasses, the maximum characteristic age for the star formation is a few hundred Myr. If one
speculates that star formation has been occurring sporadically for this timescale, then the starburst seen by
Storchi-Bergmann et al. (2005) is the most recent active episode.
In order to make a rough estimate of the supernova rate in the central region we make use of measure-
ments reported by Hummel et al. (1987). They find an unresolved component (size < 0.1′′) with 5 GHz flux
density 3.5 ± 0.3 mJy, but at lower resolution there is a 4.1 ± 0.3 mJy component of size 1′′. As discussed
in Section 3 we assume that the difference – albeit with only marginal significance – of 0.6 ± 0.4 mJy is
due to star formation in the central region, which implies a supernova rate of 6 × 10−4 yr−1 and hence
– 24 –
1010 νSN/LK ∼ 1.3, a value consistent with rather more recent star formation. Indeed, when compared to
Fig. 4, this and the low WBrγ imply a young age and short star formation timescale. For τSF = 10 Myr
the age is 60–70 Myr; for an instantaneous burst of star formation, the age would be ∼ 10 Myr, broadly
consistent with that of Storchi-Bergmann et al. (2005).
Thus, although our data do not uniquely constrain the age of the starburst in the nucleus of NGC 1097,
they do indicate that recent star formation has occurred; and they are consistent with a very young compact
starburst similar to that derived from optical and UV data.
A.2.5. NGC 1068
Evidence for a stellar core in NGC 1068 with an intrinsic size scale of ∼ 45 pc was first presented by
Thatte et al. (1997). Based on kinematics measured in large (2–4′′) apertures, they assumed the core was
virialized and estimated a mass-to-light ratio based on this assumption leading to an upper limit on the
stellar age of 1600Myr. Making a reasonable correction for an assumed old component lead to a younger
age of 500 Myr.
Stellar kinematics from optical integral field spectra (Emsellem et al. 2006a; Gerssen et al. 2006) show
evidence for a drop in the stellar velocity dispersion in the central few arcsec to σ∗ ∼ 100 km s
−1, inside a
region of higher 150–200km s−1 dispersion (presumably the bulge). Our near infrared adaptive optics data
are able to fully resolve the inner region where σ∗ drops, as shown in Fig. 21. As for NGC 1097, the velocity
distribution of the stars was derived through kinemetry, again making use of the uniformity of the stellar
velocity field to justify the simplifying assumption that the position angle and inclination do not change
significantly in the central 4′′. The derived inclination of 40◦ and position angle of 85◦ are quantitatively
similar to those found by other authors in the central few to tens of arcseconds (Emsellem et al. 2006a;
Gerssen et al. 2006; Garćıa-Lorenzo et al. 1999). The uniformity of the stellar kinematics is in contrast to
molecular gas kinematics, as traced via the 1-0 S(1) line, which are strongly perturbed and show several
distinct structures superimposed. These are too complex to permit a comparably simple analysis and will
be discussed, together with the residuals in the stellar kinematics in a future work (Mueller Sánchez et al.
in prep).
The crucial result relevant here is that at our H-band resolution of 0.10′′ we find that σ∗ reduces from
130 km s−1 at 1–2′′ to only 70 km s−1 in the very centre. That there is in the same region an excess in the
stellar continuum is demonstrated in Fig. 22. Here we show the radial profile of the stellar continuum from
both SINFONI integral field spectra out to a radius of 2′′ and NACO longslit spectra out to 5′′ (350 pc).
At radii 1–5′′, corresponding roughly to the region of high stellar dispersion measured by Emsellem et al.
(2006a), the profile is well matched by an r1/4 law, as one might expect for a bulge. At radii r < 1′′ – the
same radius at which we begin to see a discernable reduction in the stellar dispersion – the stellar continuum
increases by as much as a factor 2 above the inward extrapolation of the profile, indicating that there is extra
emission. As for NGC 1097, the combined signature of dynamically cool kinematics and excess emission is
strong evidence for a nuclear disk which has experienced recent star formation.
We can make an estimate of the characteristic age of the star formation in the central arcsec based on
the mass-to-light ratio in a similar way to Thatte et al. (1997). Because the stars appear to lie in a disk,
we estimate the dynamical mass as described in Section 3 from the stellar kinematics, using the rotation
velocity and applying a correction for the dispersion. The stellar rotation curve is essentially flat at 0.1–0.5′′,
with V∗ = 45 km s
−1 (corrected for inclination). We also take σ∗ = 70 km s
−1, which is the central value and
– 25 –
hence least biassed by the high dispersion bulge stars. These lead to a mass of 1.3× 108 M⊙ within r = 0.5
(35 pc), and a mean surface density of 3×104 M⊙ pc
−2. Correcting for the non-stellar continuum, the H-band
magnitude (which the behaviour of σ∗ indicates is dominated by the disk emission) in the same region is
11.53 mag. For H − K = 0.15 mag (Fig. 4), we find LK = 4.3 × 10
7 L⊙ and hence M/LK = 3 M⊙/L⊙.
If no star formation is on-going, this implies a characteristic age of 200–300Myr fairly independent of the
timescale (for τSF . 100 Myr, see Fig. 4) on which stars were formed. We note that this is significantly
younger than the age estimated by Thatte et al. (1997) primarily because their mass was derived using a
higher σ∗ corresponding to the bulge stars.
The assumption of no current star formation is clearly demonstrated by the Brγ map in Fig. 23. Away
from the knots of Brγ, which are associated with the coronal lines and the jet rather than possible star
formation, the equivalent width is WBrγ ∼ 4Å. This is significantly less than that for continuous star
formation of any age. Thus, while it seems likely that star formation has occurred in the last few hundred
Myr, it also seems an unavoidable conclusion that there is no current star formation.
To complete our set of diagnostics for NGC 1068, we consider also the radio continuum. This is clearly
dominated by phenomena associated with the AGN and jets, and our best estimate of the flux density away
from these features is given by the lowest contour in maps such as Figure 1 of Gallimore et al. (2004). From
this we estimate an upper limit to the 5 GHz continuum associated with star formation of 128 mJy within
r < 0.5′′. However, converting to a supernova rate and comparing to the K-band stellar luminosity yields a
limit that is not useful, being an order of magnitude above the largest expected values.
A.2.6. NGC 3783
At near infrared wavelengths, the AGN in NGC 3783 is remarkably bright. Integrated over the central
0.5′′ less than 4% of the K-band continuum is stellar. In addition, the broad Brackett lines are very strong
and dominate the H-band. Both of these phenomena are immediately clear from the H- and K-band spectra
in Fig. 24. However, it does mean that the spatial resolution can be measured easily from both the non-stellar
continuum and the broad emission lines (see Section 2). We find the K-band PSF to be symmetrical with a
FWHM of 0.17′′.
Due to the ubiquitous Brackett emission in the H-band we were unable to reliably trace the stellar
absorption features and map out the stellar continuum. Instead we have used the CO 2-0 bandhead at
2.3µm even though the dilution at the nucleus itself is extreme. The azimuthally averaged radial profile
is shown in Fig. 25 together with the PSF for reference. At radii from 0.2′′–1.6′′ (the maximum we can
measure) the profile is well fit by an r1/4 de Vaucouleurs law with Reff = 0.6
′′ (120 pc). As has been the case
previously, at smaller radii we find an excess that here is perhaps marginally resolved. Thus a substantial
fraction of the near infrared stellar continuum in the central region is likely to originate in a population of
stars distinct from the bulge.
We were unable to measure the stellar kinematics due to the limited signal-to-noise. Instead, we used
the molecular gas kinematics to estimate the dynamical mass. As before, we used kinemetry to derive the
position angle of −14◦ and the inclination in the range 35–39◦. This orientation is consistent with the
larger (20′′) scale isophotes in the J-band 2MASS image and implies that in NGC 3783 there is no significant
warp on scales of 50 pc to 4 kpc. A small inclination is also consistent with its classification as a Seyfert 1.
Adopting these values, the resulting rotation curve is shown in Fig. 26. At very small radii the rising rotation
curve may be the result of beam smearing across the nucleus. At r > 0.2′′, the falling curve suggests that
– 26 –
the rotation is dominated by the central (r < 0.2′′) mass, perhaps the supermassive black hole. We estimate
the dynamical mass within a radius of 0.3′′ (60 pc), corresponding to the point where the excess continuum
begins and also where the rotation curve appears to be unaffected by beam smearing.
Taking Vrot = 60 km s
−1 and σ = 35 km s−1 we derive a dynamical mass of Mdyn = 1.0 × 10
8 M⊙. The
black hole mass of 3 × 107 M⊙ (from reverberation mapping, Peterson et al. 2004) is only 30% of this, and
so cannot be dominating the dynamics on this scale unless its mass is underestimated. With respect to this,
we note that Peterson et al. (2004) claim the statistical uncertainty in masses derived from reverberation
mapping is about a factor 3. Alternatively, there may be a compact mass of gas and stars at r < 0.3′′.
However, including σ in the mass estimate implicitly assumes that the dispersion arises from macroscopic
motions. On the other hand, because we are observing only the hot H2, it is possible that the dispersion is
dominated by turbulence arising from shocks or UV heating of clouds that generate the 1-0 S(1) emission –
issues that are discussed in more detail by Hicks et al. (in prep.). In this case we will have overestimated
the dynamical mass. Excluding σ from the mass estimation yields Mdyn = 5 × 10
7 M⊙. We consider these
two estimates as denoting the maximum range of possible masses. Subtracting MBH then gives a mass
of stars and gas in the range (2–7)×107 M⊙, implying a mass surface density of 1700–6000M⊙ pc
−2 and
M/LK = 0.6–2.1M⊙/L⊙. Based on these ratios alone, Fig 4 indicates that the characteristic age of the star
formation may be as low as ∼ 70 Myr, although it could also be an order of magnitude greater. Without
additional diagnostics we cannot discriminate further.
We are unable to use Brγ as an additional constraint on the star formation history. Its morphology and
velocity field are similar to that of [Sivi], and rather different from the 1-0 S(1). It shows an extension to the
north which appears to be outflowing at > 50 km s−1 (Fig 27) – perhaps tracing an ionisation cone. Since
the Brγ resembles the [Sivi], it is reasonable to conclude that it too is associated with the AGN rather than
star formation. Thus the equivalent width of Brγ (with respect to the stellar continuum) of WBrγ = 30Å
represents an upper limit to that associated with star formation.
The radio continuum in the nucleus of NGC 3783 has been measured with several beam sizes at 8.5 GHz.
For a beam of 1.59′′ × 0.74′′, Morganti et al. (1999) found it was unresolved with a flux density of 8.15 ±
0.24 mJy. With a smaller ∼ 0.25′′ beam, Schmitt et al. (2001) measured a total flux density of 8.0 mJy
dominated by an unresolved component of 7.7± 0.05mJy. At smaller scales still of ∼ 0.03′′ corresponding to
6 pc, Sadler et al. (1995) placed an upper limit on the 8.5 GHz flux density of 7 mJy. Taken together, these
results imply that there is some modest 8.5 GHz radio continuum of 0.7–1 mJy extended on scales of 0.3–1′′.
Based on this we estimate a supernova rate as described in Section 3 of ∼ 0.007 yr−1, and hence a ratio
1010 νSN/LK ∼ 2. Given that the unresolved radio continuum on the smallest scales is an upper limit, the
extended component may be stronger and hence the true νSN/LK ratio may be greater than that estimated
here. Fig 4 then puts a relatively strong limit of ∼ 50 Myr on the maximum age of the star formation.
This age is fully consistent with that above associated with our lower mass estimate. The value of
WBrγ < 30Å above does not impose additional constraints, although we note that if the Brγ flux associated
with star formation is only a small fraction of the total then it would imply that the timescale over which
the star formation was active is no longer than a few times ∼ 10 Myr. Therefore in the nucleus (r < 0.3′′) of
NGC 3783 we adopt 50–70Myr as the age of the star formation and Mdyn = 2 × 10
7 M⊙ as the dynamical
mass excluding the central supermassive black hole.
– 27 –
REFERENCES
Abuter R., Schreiber J., Eisenhauer F., Ott T., Horrobin M., Gillesen S., 2006, NewAR, 50, 398
Anantharamaiah K., Viallefond F., Mohan R., Goss W., Zhao J., 2000, ApJ, 537, 613
Asari N., Vega L., Garcia-Rissmann A., González Delgado R., Storchi-Bergmann T., Cid Fernandes R., 2007,
in Galaxy Evolution Across the Hubble Time, eds Combes F., Palous J., IAU Symposium 235, 71
Baganoff F., et al., 2003, ApJ, 591, 891
Balcells M., Graham A., Domı́nguez-Palmero L., Peletier R., 2003, ApJ, 582, L79
Bender R., Burstein D., Faber S., 1992, ApJ, 399, 462
Bland-Hawthorne J., Lumsden S., Voit G., Cecil G., Weisheit J., 1997, Ap&SS, 248, 177
Bondi H., 1952, MNRAS, 112, 195
Bonnet H., et al., 2004, The ESO Messenger, 117, 17
Brand K., et al., 2006, ApJ, 663, 204
Capetti A., Axon D., Macchetto F., Marconi A., Winge C., 1999, ApJ, 516, 187
Carilli C., Wrobel J., Ulvestad J., 1998, AJ, 115, 928
Chapman S., Morris S., Alonso-Herrero A., Falcke H., 2000, MNRAS, 314, 263
Chevalier R., 1977, ARA&A, 15, 175
Chevalier R., Fransson C., 2001, ApJ, 558, L27
Cid Fernandes R., Gu Q., Melnick J., Terlevich E., Terlevich R., Kunth D., Rodrigues Lacerda R., Joguet
B., 2004, MNRAS, 355, 273
Cid Fernandes R., Heckman T., Schmitt H., González Delgado R., Storchi-Bergmann T., 2001, ApJ, 558, 81
Colina L., Alberdi A., Torrelles J., Panagia N., Wilson A., 2001, ApJ, 553, L19
Condon J., Huang Z.-P., Yin Q., Thuan T., 1991, ApJ, 378, 65
Condon J., 1992, ARA&A, 30, 575
Côté P., et al., 2006, ApJS, 165, 57
Cuadra J., Nayakshin S., Springel V., Di Matteo T., 2006, MNRAS, 366, 358
Dasyra K., et al., 2006, ApJ, 651, 835
Davies R., Sternberg A., Lehnert M., Tacconi-Garman L., 2003, ApJ, 597, 907
Davies R., Tacconi L., Genzel R., 2004a, ApJ, 602, 148
Davies R., Tacconi L., Genzel R., 2004b, ApJ, 613, 781
Davies R., Sternberg A., Lehnert M., Tacconi-Garman L., 2005, ApJ, 633, 105
– 28 –
Davies R., et al., 2006, ApJ, 646, 754
Davies R., 2007a, MNRAS, 375, 1099
Davies R., 2007b in The 2007 ESO Instrument Calibration Workshop, ESO Astrophysics Symposia,
arXiv:astro-ph/0703044
Downes D., Solomon P., 1998, ApJ, 507, 615
Eisenhauer F., et al., 2003, in Instrument Design and Performance for Optical/Infrared Ground-based Tele-
scopes, eds. Masanori I., Moorwood A., Proc. SPIE, 4841, 1548
Emsellem E., Greusard D., Combes F., Friedli D., Leon S., Pécontal E., Wozniak H., 2001, A&A, 368, 52
Emsellem E., Fathi K., Wozniak H., Ferruit P., Mundell C., Schinnerer E., 2006a, MNRAS, 365, 367
Emsellem E., 2006b, in The Central Engine of Active Galactic Nuclei, eds Ho L., Wang J.-M. (San Francisco:
Emsellem E., 2006c, in Mapping the Galaxy and Nearby Galaxies, eds Wada K., Combes F. (Springer)
Fathi K., Storchi-Bergmann T., Riffel R., Winge C., Axon D., Robinson A., Apetti A., Marconi A., 2006,
ApJ, 641, L25
Förster Schreiber N., Genzel R., Lutz D., Sternberg A., 2003, ApJ, 599, 193
Förster Schreiber N., 2000, AJ, 120, 2089
Gallimore J., Baum S., O’Dea C., 2004, ApJ, 613, 794
Garćıa-Lorenzo B., Mediavilla E., Arribas S., 1999, ApJ, 518, 190
Genzel R., Weitzel L., Tacconi-Garman L.E., Blietz M., Cameron M., Krabbe A., Lutz D., Sternberg A.,
1995, ApJ, 444, 129
Genzel R., et al., 2003, ApJ, 594, 812
Gerssen J., Allington-Smith J., Miller B., Turner J., Walker A., 2006, MNRAS, 365, 29
Glass I., 1984, MNRAS, 211, 461
González Delgado R., Heckman T., Leitherer C., 2001, ApJ, 546, 845
González Delgado R., Cid Fernandes R., 2005, in Starbursts: from 30 Doradus to Lyman Break Galaxies,
eds de Grijs R., González Delgado R., (Dordrecht: Springer), ASSL, 329, 263
Graham A., Guzmán R., 2003, AJ, 125, 2936
Granato G., de Zotti G., Silva L., Bressan A., Danese L., 2004, ApJ, 600, 580
Greenhill L., et al., 2003, ApJ, 590, 162
Gu Q., Dultzin-Hacyan D., de Diego J., 2001, RMxAA, 37, 3
Heckman T., González-Delgado R., Leitherer C., Meurer G., Krolik J., Wilson A., Koratkar A., Kinney A.,
1997, ApJ, 482, 114
http://arxiv.org/abs/astro-ph/0703044
– 29 –
Huang Z., Thuan T., Chevalier R., 1994, ApJ, 424, 114
Hummel E., van der Hulst J., Keel W., 1987, A&A, 172, 32
Ivanov V., Rieke G., Groppi C., Alonso-Herrero A., Rieke M., Engelbracht C., 2000, ApJ, 545, 190
Johuet B., Kunth D., Melnick J., Terlevich R., Terlevich E., 2001, A&A, 380, 19
Kennicutt Jr R., 1998, ApJ, 498, 541
Knapen J., 2004, in Penetrating bars through masks of cosmic dust : the Hubble tuning fork strikes a new note,
eds Block D., Puerari I., Freeman K., Groess R., Block E. (Dordrecht: Kluwer Academic Publishers),
ASSL, vol. 319, p.189
Kormendy J., Sanders D., 1992, ApJ, 390, L53
Kormendy J., Fisher D., Cornell M., Bender R., 2007, ApJ, submitted
Krajnović D., Cappellari M., de Zeeuw T., Copin Y., 2006, MNRAS, 366, 787
Krolik J., 2007, ApJ, in press
Krumholz M., McKee C., Klein R., 2006, ApJ, 638, 369
Leitherer C., et al., 1999, ApJS, 123, 3
Lenzen R., et al., 2003, in Instrument Design and Performance for Optical/Infrared Ground-based Telescopes,
eds. Masanori I., Moorwood A., Proc. SPIE, 4841, 944
Lewis K., Eracleous M., 2006, ApJ, 642, 711
Lodato G., Bertin G., 2003, A&A, 398, 517
Lonsdale C., Smith H., Lonsdale C., 1995, ApJ, 438, 632
Machalski J., Chyży K., Stawarz  L., Kozie lD., 2007, A&A, 462, 43
Maiolino R., Krabbe A., Thatte N., Genzel R., 1998, ApJ, 493, 650
Martinez-Sansigre A., Rawlings S., Lacy M., Fadda D., Jarvis M., Marleau F., Simpson C., Willott C., 2006,
MNRAS, 370, 1479
Martini P., Regan M., Mulchaey J., Pogge R., 2003, ApJ, 589, 774
Márquez I., Masegosa J., Durret F., González Delgado R., Moles M., Maza J., Pérez E., Roth M., 2003,
A&A, 409, 459
Melia F., ApJ, 387, L25
Mihos C., Hernquist L., 1994, ApJ, 437, L47
Morganti R., Tsvetanov Z., Gallimore J., Allen M., 1999, A&AS, 137, 457
Mueller Sánchez F., Davies R., Eisenhauer F., Tacconi L., Genzel R., Sternberg A., 2006, A&A, 454, 481
Nelson C., Whittle M., 1995, ApJS, 99, 67
– 30 –
Nikolaev S., Weinberg M., 1997, ApJ, 487, 885
Oliva E., Origlia L., Kotilainen J., Moorwood A., 1995, A&A, 301, 55
Origlia L., Moorwood A., Oliva E., 2003, A&A, 280, 536
Ozernoy L, Genzel R., 1996, eds Blitz L., Teuben P., (Kluwer) Proc. IAU Symposium 169, p.181
Ozernoy L., Genzel R., Usov V., 1997, MNRAS, 288, 1997
Paumard T., Maillard J.-P., Morris M., Rigaut F., 2001, A&A, 366, 466
Paumard T., et al., 2006, ApJ, 643, 1011
Peterson B., et al., 2004, ApJ, 613, 682
Pier E., Krolik J., 1992, ApJ, 399, L23
Pier E., Antonucci R., Hurt T., Kriss G., Krolik J., 1994, ApJ, 428, 124
Prieto A., Maciejewski W., Reunanen J., 2005, AJ, 130, 1472
Rigby J., Rieke G., Donley J., Alonso-Herrero A., Pérez-González P., 2006, ApJ, 645, 115
Rousset G., et al., 2003, in Adaptive Optical System Technologies II, eds. Wizinowich P., Bonaccini D., Proc.
SPIE, 4839, 140
Sadler E., Slee O., Reynolds J., Roy A., 1995, MNRAS, 276, 1273
Sarzi M., Shields J., Pogge R., Martini P., 2007, in Stellar Populations as Building Blocks of Galaxies, eds
Peletier R., Vazdekis A., IAU Symposium 241
Schmitt H., Ulvestad J., Antonuccci R., Kinney A., 2001, ApJS, 132, 199
Scoville N., et al., 2000, AJ, 119, 991
Scoville N., 2003, JKAS, 36, 167
Shapiro K., Gerssen J., van der Marel R., 2003, AJ, 126, 2707
Smith H., Lonsdale C., Lonsdale C., Diamond P., 1998, ApJ, 493, L17
Springel V., Di Matteo T., Hernquist L., 2005, MNRAS, 361, 776
Sternberg A., 1998, ApJ, 506, 721
Sternberg A., Hoffman T., Pauldrach A., 2003, ApJ, 599, 1333
Storchi-Bergmann T., González Delgado R., Schmitt H., Cid Fernandes R., Heckman T., 2001, ApJ, 559,
Storchi-Bergmann T., et al., 2003 ApJ, 598, 956
Storchi-Bergamnn T., Nemmen R., Spinelli P., Eracleous M., Wilson A., Filippenko A., Livio M., 2005, ApJ,
624, L13
Thatte N., Quirrenbach A., Genzel R., Maiolino R., Tecza M., 1997, ApJ, 490, 238
– 31 –
Thompson T., Quataert E., Murray N., 2005, ApJ, 630, 167
Vazquez G., Leitherer C., 2005, ApJ, 621, 695
Wehrle A., Morris M., 1988, AJ, 95, 1689
Winters J., Le Bertre T., Jeong K., Nyman L.-Å., Epchtein N., 2003, A&A, 409, 715
Woo J.-H., Urry C., 2002, ApJ, 579, 530
Wozniak H., Combes F., Emsellem E., Friedli D., 2003, A&A, 409, 469
This preprint was prepared with the AAS LATEX macros v5.2.
– 32 –
Table 1. Table of Observations
Object Banda Res.b (′′) Date Instrument
Mkn 231c H 0.176 May ’02 Keck, NIRC2
NGC 7469c K 0.085 Nov ’02 Keck, NIRSPAO
K 0.15 Jul ’04 VLT, SINFONI
Circinusc K 0.22 Jul ’04 VLT, SINFONI
NGC 3227c K 0.085 Dec ’04 VLT, SINFONI
IRAS 05189-2524 H 0.12 Dec ’02 VLT, NACO
NGC 2992 K 0.30 Mar ’05 VLT, SINFONI
NGC 1097 H 0.245 Oct ’05 VLT, SINFONI
NGC 1068 H 0.10 Oct ’05 VLT, SINFONI
H 0.13 Dec ’02 VLT, NACO
NGC 3783 K 0.17 Mar ’05 VLT, SINFONI
aBand used for determining the quantitative star formation proper-
ties. NGC 1097, NGC 1068, and NGC 3783 were actually observed in H
and K bands.
bSpatial resolution (FWHM) estimated from the data itself, using the
methods described in Section 2.
cReferences to detailed studies of individual objects: Mkn 231
(Davies et al. 2004b), NGC 7469 (Davies et al. 2004a), Circinus,
(Mueller Sánchez et al. 2006), NGC 3227 (Davies et al. 2006).
– 33 –
Table 2. Summary of basic data for AGN
Object Classificationa Distance logLbol
b logMBH
Ref.c
Mpc for MBH
Mkn 231 ULIRG, Sy 1, QSO 170 12.5 7.2 1
NGC 7469 Sy 1 66 11.5 7.0 2
Circinus Sy 2 4 10.2 6.2 3
NGC 3227 Sy 1 17 10.2 7.3 4
IRAS 05189-2524 ULIRG, Sy 1 170 12.1 7.5 1
NGC 2992 Sy 1 33 10.7 7.7 5
NGC 1097 LINER, Sy 1 18 10.9 8.1 6
NGC 1068 Sy 2 14 11.5 6.9 7
NGC 3783 Sy 1 42 10.8 7.5 2
a Classifications are taken primarily from the NASA/IPAC Extragalactic Database. In
addition, we have labelled as Seyfert 1 those for which we have observed broad (i.e. FWHM
> 1000 km s−1) Brγ; see also Fig 1.
b Calculated in the range 8–1000µm from the IRAS 12–100µm flux densities; with an
additional correction for optical and near-infrared luminosity in cases where appropriate.
c References for black hole masses: (1) Dasyra et al. (2006); (2) Peterson et al.
(2004); (3) Greenhill et al. (2003); (4) Davies et al. (2006); (5) Woo & Urry (2002); (6)
Lewis & Eracleous (2006); (7) Lodato & Bertin (2003)
Table 3. Measured & Derived Properties of the Nucleia
Object radius logLK∗
Σdyn WBrγ Mdyn/LK 10
10νSN/LK age 〈SFR〉
′′ pc 104M⊙ pc
−2 Å M⊙/L⊙ yr
−1L−1
Myr M⊙ yr
−1 kpc−2
Mkn 231b 0.6 480 9.3 9.8 0.9 — 3.1 20 120–250 25–50
NGC 7469 0.4 128 8.5 8.7 1.0 11 1.6 3 110–190 50–100
Circinus 0.4 8 6.2 7.5 17 30 23 1.5 80 ∼ 70
NGC 3227c 0.4 32 7.8 8.0 3.7 4 1.9 2.2 40 ∼ 380
IRAS 05189-2524 0.55 450 9.3 — — 4 — 5 50–100 30–70
NGC 2992d 0.4 64 7.5 — — <12 — 1 — —
NGC 1097e 0.25 22 6.7 8.2 1.3 1 4.5 1.4 8 ∼ 80
NGC 1068 0.5 35 7.6 8.1 3.4 4 3.0 <20 200–300 90–170
NGC 3783f 0.3 60 7.5 7.3 0.2 <30 0.6 2 50–70 30–60
a The methods used to measure these quantities (within the radii given) are described in Section 3. Specific issues associated with
individual objects are discussed in Appendix A.
b Mdyn depends strongly on even small changes to the inclination; here it is given for i = 10
◦. Correcting Mdyn for an estimate of the gas
mass given in Downes & Solomon (1998) yields M/LK = 2.3M⊙/L⊙.
c The best star formation models indicate that M/LK is much less than the limit given here using the dynamical mass.
d It is likely that much of the narrow Brγ in the nuclear region here is associated with an ionisation cone. In addition, the high stellar
velocity dispersion, even on the smallest scales we have been able to measure, suggests that the K-band light is dominated by the bulge.
e Correcting WBrγ for the old stellar population would probably yield a value in the range 2–5Å. Even on this scale the dynamical mass
is dominated by the supermassive black hole. Both Σdyn and M/LK are estimated after subtracting MBH.
f Much of the Brγ here is outflowing and hence associated with the AGN. Mdyn is derived from gas kinematics as described in the text.
Both Σdyn and M/LK are estimated after subtracting MBH.
– 35 –
Fig. 1.— K-band spectra showing broad Brγ emission in 3 AGN which are not usually classified as Seyfert 1.
Top: NGC 3227 (0.25′′ aperture); Middle: IRAS 05189−2524 (1′′ aperture); Bottom: NGC 2992 (0.5′′ aper-
ture). The most prominent emission and absorption features are marked.
Fig. 2.— Left: equivalent width of the CO 2-0 and CO 6-3 features for various stars and galaxies. The late-
type supergiant stars (skeletal star shapes) and giant stars (open star shapes) are taken from Origlia et al.
(1993). The galaxies (denoted ‘E’ for elliptical, ‘S’ for spiral, and ‘H’ for star forming Hii galaxy) are from
Oliva et al. (1995). The dashed box encloses the region for which there is no more than 20% deviation
from each of the values WCO2−0 = 12Å and WCO6−3 = 4.5Å. Centre and right: calculated WCO6−3 and
WCO2−0 respectively from STARS for several different star formation histories. Each line is truncated when
the cluster luminosity falls below 1/15 of its maximum. In each case, the dashed lines show typical values
adopted, and the dotted lines denote a range of ±20%.
– 36 –
Fig. 3.— Left: H-K colour of star clusters with different star formation timescales and ages, as calculated by
STARS. Right: Ratio of bolometric to K-band luminosity. Although the range of 20–200 initially appears
large, the uncertainty on an intermediate value of 60 is only 0.3 dex. This is small compared to the range
of interest in the paper, which is several orders of magnitude.
Fig. 4.— Various diagnostics calcuated with STARS for several star formation timescales, as functions of
age: Brγ equivalent width, supernova rate, and mass-to-light ratio. Note that all are normalised to the
K-band stellar continuum; and that LK is the total luminosity in the 1.9–2.5µm band in units of bolometric
solar luminosity (1Lbol = 3.8 × 10
26 W ), rather than the other frequently used monochromatic definition
which has units of the solar K-band luminosity density.
– 37 –
Fig. 5.— Size scales of nuclear star forming regions. The profile has been determined from the CO absorption
features in the H or K band, which are approximately independent of star formation history (see text for
details). For longslit data, spatial profiles have been averaged; for integral field data, azimuthally averaged
proifles are shown. For NGC1068, data at two different pixel scales are shown (corresponding to the solid
and dashed brown lines). The horizontal dotted line is drawn at half-maximum height, to assist in estimating
size scales by eye.
– 38 –
Fig. 6.— A STARS star formation model based on the main characteristics of the observed starbursts,
which is illustrative of a ‘typical’ nuclear starburst that we have observed. The scaling is fixed as 2× 109 M⊙
(typical of that within a 30 pc radius; Fig. 9) at an age of 100 Myr (the typical age in Table 3). The star
formation timescale is τSF = 10 Myr to reproduce a low Brγ equivalent width. The panels are, from top left:
(a) star formation rate (SFR); (b) number of ionising photons (QLyc, proportional to the Brγ luminosity);
(c) supernova rate (SNR); (d) bolometric luminosity (Lbol); (e) K-band luminosity (LK); (f) cumulative
number of supernovae; (g) cumulative mass that has been recycled back into the ISM by supernovae and
winds; (h) mass loss rate from stars; (i) rate at which the lost mass can in principle be accreted onto a central
supermassive black hole (due to its outflow speed; see text for details). In the last two panels, the mass loss
is split into that due to OB and Wolf-Rayet stars and supernovae (dotted lines), and that due to late-type
and AGB stars (dashed lines). Stellar mass loss in STARS is accounted for at the end of each star’s life as
the difference in mass between the original (ZAMS) stellar mass and the remnant mass as the end-product
of its stellar evolution, as described in Sternberg (1998). For this reason, the mass loss rates from OB and
Wolf-Rayet stars do not appear explicitly in panel (h) at very young cluster ages.
– 39 –
Fig. 7.— Mean enclosed mass surface density as a function of radius. The points are from data given
in Table 3; the dashed lines represent the mass models derived for NGC 7469 and Mkn 231 (Davies et al.
2004a,b). The galaxies all follow the same trend towards increasing densities in the central regions.
Fig. 8.— Integrated bolometric luminosity of the young stars Lbol∗ as a fraction of that of the galaxy
Lbol−galaxy, plotted as a function of radius. Lbol∗ is calculated from the stellar LK or LH assuming that,
on the generally small scales here, all the near infrared stellar continuum originates in the young stars. On
10 pc scales the contribution of young stars is at most a few percent of the galaxy’s total luminosity, while
on kpc scales it may be significant and hence comparable to the AGN luminosity. For NGC 1068, data at
two different pixel scales are shown (corresponding to the solid and dashed brown lines).
– 40 –
Fig. 9.— Integrated bolometric luminosity of young stars Lbol∗ (see Fig. 8) as a function of radius. For
comparison, the dotted line has constant surface brightness. For NGC 1068, data at two different pixel scales
are shown (corresponding to the solid and dashed brown lines).
Fig. 10.— Luminosity to mass ratio calculated by STARS as a function of age for different star formation
timescales. The dotted lines show how the ratio would vary if gas were fed in to a cluster at the same rate
as it was converted into stars. The solid lines assume that the gas is present at the start, but at the end has
all been processed in stars. A cluster can only exceed 500 L⊙/M⊙ for a timescale of ∼ 10 Myr.
– 41 –
Fig. 11.— Graph showing how the luminosity of an AGN might be related to the age of the most recent
episode of nuclear star formation. On the left is shown the luminosity in solar units; on the right, it is with
respect to the Eddington luminosity for the black hole. Generally the luminosity of the AGN is not well
known and so we have approximated it by 0.5Lbol, and adopted an uncertainty of a factor 2. The starburst
age refers to our best estimate of the most recent episode of star formation within the central 10–100pc, as
given in Table 3. See the text for details of the adopted uncertainties.
– 42 –
Fig. 12.— Spectrum (in grey) of IRAS 05189-2524, integrated over two 0.22′′ wide sections centered ±0.27′′
either side of the nucleus, and which have been shifted to match their velocities. Overplotted in black is
a fit to the continuum, comprising spectra of various supergiant stars. The main absorption and emission
features have been identified.
Fig. 13.— Spatial profiles of non-stellar (left), stellar (centre), and total (right) continuum for IRAS 05189-
2524 (1′′ = 800 pc). The first two have been derived at each point along the spatial extent of the slit from
the spectral slope and the stellar absorption features respectively. A comparison of their sum (dashed line
in right panel) to the total continuum indicates that the decomposition appears to be reasonable.
– 43 –
Fig. 14.— Maps of NGC 2992 (1′′ = 160 pc). From left to right: stellar continuum, narrow Brγ, Brγ velocity
(−150 to +150 km s−1), and Brγ dispersion (0 to 200 km s−1). For reference, on each panel are superimposed
contours of the Brγ flux (8, 16, 32, and 64% of the peak). The symbol plotted on the map of the line flux
indicates the centre of the broad Brγ and non-stellar emission. The narrow Brγ emission extends far more
to the north west than the stellar continuum. And, particularly on the south western edge, it exhibits a
blue shifted velocity and high dispersion. All these are consistent with an interpretation as the apex of an
ionisation cone.
Fig. 15.— Radial profiles of velocity and dispersion for the stars in NGC 2992 (1′′ = 160 pc). The 2D maps
were analysed using the kinemetric technique described by Krajnović et al. (2006) which yielded a position
angle of 24◦ and an inclination of ∼ 40◦, not dissimilar to the isophotal values of 30◦ and 50◦ respectively.
The rotation curve has been corrected for the inclination.
– 44 –
Fig. 16.— Radial profile of the stellar continuum in NGC 2992 (1′′ = 160 pc), derived from isophotal analysis.
Solid circles denote the stellar continuum (i.e. already corrected for the non-stellar component). Overplotted
with triangles are an r1/4 law (top panel) and an exponential profile (bottom panel). The profiles were fitted
at radii r > 0.5′′ and extrapolated inwards, convolved with the PSF which is shown as open squares. Both
fits are equally good at r > 0.5′′, but only the exponential suggests there might be excess continuum at the
centre, arising from a distinct stellar population. This is therefore inconclusive.
– 45 –
Fig. 17.— Spectrum of NGC 1097 (thick grey line), extracted within an aperture of radius 0.2′′ and scaled
arbitrarily. Overdrawn (thin black line) is a match to the stellar continuum constructed from template
spectra of several late-type supergiant stars and a blackbody function representing the non-stellar component.
Notably, Brγ in the nucleus is extremely weak even in the nucleus.
Fig. 18.— Radial profile of the stellar continuum in NGC 1097 (1′′ = 80 pc). The solid circles denote the
stellar continuum (i.e. already corrected for the non-stellar component). The triangles denote an r1/4 profile
fitted to radii r > 0.5′′ and extrapolated inwards. This model has been convolved with the PSF, shown as
open squares for comparison. Note that even though an exponential profile might match the data equally
well, an r1/4 profile provides a stronger constraint on whether there is excess continuum at the centre.
– 46 –
Fig. 19.— Radial profiles of velocity and dispersion for the gas and stars in NGC 1097 (1′′ = 80 pc). The 2D
maps were created by convolving template spectra (i.e unresolved line profile for the gas, stellar template
for the stars) with a Gaussian and minimising the difference with respect to the galaxy spectrum at each
spatial pixel. These were then analysed using the kinemetric technique described by Krajnović et al. (2006)
which yielded the same position angle of −49◦ for the gas and stars, and similar inclinations of 32◦ and 43◦
respectively. The rotation curve has been corrected for the inclination.
Fig. 20.— Maps of K-band non-stellar continuum (left), stellar continuum (centre), and Brγ line flux (right)
for the central few arcsec of NGC 1097 (1′′ = 80 pc). In each case, the centre (as defined by the non-stellar
continuum) is marked by a crossed circle. The colour scale is shown on the right, as percentage of the peak
in each map.
– 47 –
Fig. 21.— Radial profiles of velocity and dispersion for stars in NGC 1068 (1′′ = 70 pc). The 2D maps were
then analysed using kinemetry Krajnović et al. (2006), yielding an inclination of 40◦ and a position angle of
85◦. The rotation curve has been corrected for the inclination. Also plotted for comparison are azimuthally
averaged radial profiles of the H-band stellar luminosity and the PSF.
Fig. 22.— Radial profiles of the H-band stellar luminosity in NGC 1068 (1′′= 70 pc) from NACO longslit
data at position angle 0◦ and 90◦ (green and red symbols; squares and triangles deonte opposite sides of the
nucleus). To trace the profile to larger radius, open symbols are each the mean of 9 points. Superimposed
are SINFONI data from Fig. 21 (blue crosses, flux scaled to match). The dashed line denotes an r1/4 law
with Reff = 1.5
′′ to match the outer profile. At r < 1′′, the stellar continuum reveals an excess above the
inward extrapolation of this profile.
– 48 –
Fig. 23.— Maps of the central few arcsec of NGC 1068 (1′′ = 70 pc): H-band non-stellar continuum (far left)
and stellar continuum (centre left); also Brγ line flux (center right) and Brγ equivalent width (far right).
In each case, the centre (as defined by the non-stellar continuum) is marked by a crossed circle. The colour
scale is shown on the right, as percentage of the peak in each map (and also as WBrγ in Å).
– 49 –
Fig. 24.— H- and K-band spectra of the central 1′′ of NGC 3783, with the prominent emission and absorption
features labelled. In the H-band it is challenging to measure the stellar absorption due to the very strong
brackett emission from the AGN’s broad line region. Instead we have used the K-band CO 2-0 bandhead
even though the dilution at this wavelength is extreme.
– 50 –
Fig. 25.— Radial profile of the stellar continuum in NGC 3783 (1′′ = 200 pc). The solid circles denote the
stellar continuum (i.e. already corrected for the non-stellar component). The triangles denote an r1/4 profile
fitted to radii 0.2 < r < 1.6′′ and extrapolated inwards. This model has been convolved with the PSF, shown
as open squares for comparison. Note that even though an exponential profile might match the data equally
well, an r1/4 profile provides a stronger constraint on whether there is excess continuum at the centre.
Fig. 26.— Rotation curve derived from the the H2 1-0 S(1) velocity field in NGC 3783 (1
′′ = 200 pc).
Also shown is the dispersion as a function of radius. The velocity field was analysed using kinemetry
(Krajnović et al. 2006) which yielded a major axis of about −14◦ and an inclination in the range 35–39◦.
The drop in velocity at r < 0.15′′ maybe due to beam smearing across the nucleus.
– 51 –
Fig. 27.— Images of the central 2′′ of NGC 3783 (1′′ = 200 pc). Top row from left: non-stellar continuum,
H2 1-0 S(1) line flux, H2 1-0 S(1) velocity. Bottom row from left: narrow Brγ line flux, [Sivi] line flux, [Sivi]
velocity. The Brγ velocity field is similar to that of [Sivi] and shows an outflow of >50 km s−1 to the north.
This is in contrast to the 1-0 S(1) velocity field which traces rotation.
	Introduction
	Sample, Observations, Data Processing
	Sample Selection
	Observations & Reduction
	PSF Estimation
	Emission/Absorption Line Characterisation
	Quantifying the Star Formation
	Isolating the stellar continuum
	Stellar colour and luminosity
	Specific Star Formation Diagnostics
	Br equivalent width
	Supernova rate
	Mass-to-light ratio
	Properties of Nuclear Star Formation
	Starburst-AGN Connection
	Conclusions
	Individual Objects
	Summary of Star Forming Properties of Galaxies already Studied
	Mkn 231
	Circinus
	NGC 3227
	Star Forming Properties of Galaxies with New Data
	NGC 7469
	IRAS 05189-2524
	NGC 2992
	NGC 1097
	NGC 1068
	NGC 3783
