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Astronomy & Astrophysics manuscript no. thickbulge˙v22 c© ESO 2014
October 14, 2014
Evidence for the concurrent growth of thick discs and central mass
concentrations from S4G imaging
S. Comero´n1,2, B. G. Elmegreen3, H. Salo1, E. Laurikainen1,2, B. W. Holwerda4, and J. H. Knapen5,6
1 University of Oulu, Astronomy Division, Department of Physics, P.O. Box 3000, FIN-90014, Finland
e-mail: seb.comeron@gmail.com
2 Finnish Centre of Astronomy with ESO (FINCA), University of Turku, Va¨isa¨la¨ntie 20, FI-21500, Piikkio¨, Finland
3 IBM Research Division, T.J. Watson Research Center, Yorktown Hts., NY 10598, USA
4 European Space Agency Research Fellow (ESTEC), Keplerlaan, 1, 2200 AG Noordwijk, The Netherlands
5 Instituto de Astrofı´sica de Canarias, E-38200 La Laguna, Tenerife, Spain
6 Departamento de Astrofı´sica, Universidad de La Laguna, E-38205 La Laguna, Tenerife, Spain
Preprint online version: October 14, 2014
ABSTRACT
We have produced 3.6µm + 4.5µm vertically integrated radial luminosity profiles of 69 edge-on galaxies from the Spitzer Survey
of Stellar Structure in Galaxies (S4G). We decomposed the luminosity profiles into a disc and a central mass concentration (CMC).
These fits, combined with thin/thick disc decompositions from our previous studies, allow us to estimate the masses of the CMCs,
the thick discs, and the thin discs (MCMC, MT, and MT). We obtained atomic disc masses (Mg) from the literature. We then consider
the CMC and the thick disc to be dynamically hot components and the thin disc and the gas disc to be dynamically cold components.
We find that the ratio between the mass of the hot components and that of the cold components, (MCMC +MT)/(Mt +Mg), does not
depend on the total galaxy mass as described by circular velocities (vc). We also find that the higher the vc, the more concentrated the
hot component of a galaxy. We suggest that our results are compatible with having CMCs and thick discs built in a short and early
high star forming intensity phase. These components were born thick because of the large scale height of the turbulent gas disc in
which they originated. Our results indicate that the ratio between the star forming rate in the former phase and that of the formation
of the thin disc is of the order of 10, but the value depends on the duration of the high star forming intensity phase.
Key words. Galaxies: evolution – Galaxies: bulges – Galaxies: spiral – Galaxies: statistics
1. Introduction
Historically, the surface brightness profiles of edge-on galax-
ies perpendicular to their mid-planes have been fitted with a
sech2 (z/z0) function, where z0 is the scale height (Spitzer 1942;
van der Kruit & Searle 1981). The roughly exponential excesses
of light discovered by Burstein (1979) and Tsikoudi (1979)
at large heights and low surface brightnesses are now known
as thick discs. They have been found in all edge-on galax-
ies where they have been sought (Dalcanton & Bernstein 2002;
Yoachim & Dalcanton 2006; Comero´n et al. 2011c). Two thick
discs have been found in at least one galaxy (Comero´n et al.
2011b).
Three main models have been proposed to explain the for-
mation of thick discs. In the first, they have a secular origin
caused by the vertical heating and/or radial migration of stars
(e.g., Villumsen 1985; Scho¨nrich & Binney 2009). However,
some studies show that radial migration does not heat the disc
(Minchev et al. 2012; Vera-Ciro et al. 2014). In the second thick
disc formation mechanism, they were thick from the beginning
as a result of a high velocity dispersion of the interstellar medium
at high redshift; the thin disc accreted later with a lower dis-
persion (e.g., Brook et al. 2004; Elmegreen & Elmegreen 2006;
Bournaud et al. 2009). In the third model, they arise from inter-
actions with satellite galaxies, either because of dynamical heat-
ing (e.g., Quinn et al. 1993; Qu et al. 2011) or because of accre-
tion of stars (e.g., Abadi et al. 2003). Whether this last family of
mechanisms is secular or not depends on the merging history of
the galaxies.
In studies of the Milky Way thick disc, some authors favour
a secular formation (e.g., Bovy et al. 2012), while others favour
an early origin (e.g., Gilmore et al. 2002; Micali et al. 2013;
Bensby et al. 2014). Extragalactic studies based on photomet-
ric decompositions have been interpreted as evidence of an early
origin (Yoachim & Dalcanton 2006; Comero´n et al. 2011a), al-
though it has also been suggested that the data for high-mass
galaxies (circular velocity vc & 120 km s−1, which corresponds
to a baryonic mass M ∼ 1010 M⊙ using the Tully-Fisher re-
lation in Zaritsky et al. 2014), are compatible with an inter-
nal secular origin (Comero´n et al. 2012, hereafter CO12). The
vc = 120 km s−1 limit is found to distinguish two groups of
disc galaxies with different properties. Low-mass galaxies have
thick discs with masses similar to those of their thin discs
(Yoachim & Dalcanton 2006; Comero´n et al. 2011a). Higher
mass galaxies (vc > 120 km s−1) usually have thick discs that
are significantly less massive than thin discs.
Central mass concentrations (CMCs) are divided into clas-
sical bulges and pseudobulges. The distinction between the two
kinds of CMCs is made using a variety of criteria (Se´rsic in-
dex, stellar populations, the presence of substructures like rings
or spiral arms, stellar kinematics, etc.), which often, but not al-
ways, provide an unambiguous classification. The picture is fur-
ther complicated by CMCs that are the superposition of a clas-
sical bulge and a pseudobulge (Erwin 2008). Classical bulges
1
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
are believed to have formed violently early in a galaxy’s his-
tory and pseudobulges are thought to have formed secularly
(e.g., Kormendy & Kennicutt 2004). However, recent numer-
ical models have shown that some pseudobulges might also
have formed quickly at high redshift (Scannapieco et al. 2011;
Inoue & Saitoh 2012, 2014). In these studies, the distinction be-
tween a classical bulge and a pseudobulge is made using the
Se´rsic index.
Here, we consider classical bulges and pseudobulges to be
part of the CMC, and we do not consider boxy/peanut inner
structures to be part of the CMC; indeed, they are considered
by some authors to be edge-on bars (Kuijken & Merrifield 1995;
Lu¨tticke et al. 2000), so we consider them as a disc substructure.
In this paper, we explore the links between the masses of
the gas, thin, and the thick discs and the CMCs in a sample of
edge-on galaxies from the Spitzer Survey of Stellar Structure
in Galaxies (S4G; Sheth et al. 2010). The S4G is a deep 3.6µm
and 4.5µm survey made with the Infrared Array Camera (IRAC;
Fazio et al. 2004) of 2352 nearby (vradio < 3000 km s−1), bright
(mB,corr < 15.5), and large (D > 1′) galaxies away from the
Galactic plane (|b| > 30o). The S4G data is public and can
be downloaded from the NASA/IPAC Infrared Science Archive
(IRSA) website1.
The paper is structured as follows. In Sect. 2 we describe our
sample, in Sect. 3 we explain the data processing, and in Sect. 4
we give our results. In Sect. 5 we discuss and interpret our results
and present our conclusions in Sect. 6. Appendix A presents the
fits made with the procedure that is described in Sect. 3.
2. The sample
The galaxy sample is based on the sample used in CO12. CO12
studied 70 nearly edge-on galaxies with Hubble types −2 ≤ T ≤
7, distances 11 Mpc ≤ d ≤ 75 Mpc (redshift-independent dis-
tances from NED, or if not available, redshift distances corrected
for the Virgo and Shapley Clusters and the Great Attractor effects
and with a Hubble-Lemaıˆtre constant H0 = 73 km s−1 Mpc−1),
and brightnesses −17.4 ≤ M3.6µm(AB) ≤ −22.2 (from the S4G
Pipeline 3; Mun˜oz-Mateos et al., in prep.).
In CO12 we performed thin/thick disc photometric decom-
positions. To ensure good fits, the galaxies were selected to have
a high edge-on surface brightness, to be undistorted, and to have
thin discs that could be vertically resolved. Additionally, we se-
lected galaxies where the disc dominated in at least the outer
∼ 50% of the optical radius, so clean thin/thick disc decomposi-
tions could be obtained. Here, compared to the CO12 sample, we
removed NGC 5084 for which we now think the bulge has a large
influence on the thin/thick disc fits and is causing the antitrunca-
tion in the radial luminosity profile (therefore, the antitruncation
in NGC 5084 seen in CO12 is not a true disc feature). This re-
duces the sample size to 69 galaxies. Our selection criteria (and
also those from the S4G sample; Sheth et al. 2010) bias the sam-
ple against early-type galaxies. Only five galaxies out of 69 have
T < 0 according to HyperLeda (Paturel et al. 2003). The prop-
erties of the galaxies in the sample can be found in Table 1.
The thin/thick disc photometric decompositions in CO12
give the relative masses of the thin and the thick discs. This is
later used to obtain the thin and thick disc masses.
0 20 40 60 80 100
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 522
-100 -50 0 50 100
-20
-10
0
10
20
Fig. 1. Radial luminosity profile of NGC 522 (yellow symbols)
and fits to it. The vertical red bars indicate the limits of the fit in
CO12. The dotted lines indicate the fitted profile and the break
radii in CO12. The continuous lines indicate the fitted profile
and the break radii in this paper. The dashed profiles indicate the
CMC and the disc contributions. The units in the inset axes are
arcseconds. The inset image is unmasked and covers the [−zu,zu]
vertical range.
3. Fits of the radial luminosity profiles
In CO12, we obtained radial luminosity profiles parallel to the
mid-plane of the galaxies for the thin and thick discs and for
the whole disc using the averages of the 3.6µm and 4.5µm im-
ages. The profiles for the whole disc were averaged over (−zu,zu),
where at zu the surface brightness averaged over some radial ex-
tent is µ(AB) = 26 mag arcsec−2. This radial extent was usually
between 0.2r25 and 0.8r25, but see CO12 for details. Here, we
only use the luminosity profiles that correspond to the whole
disc.
To produce the luminosity profiles we masked the fore-
ground stars and the background galaxies. We expanded the
masks in CO12 so the masking now also covers the region with
R < 0.2r25.
The luminosity profiles in CO12 were fitted with a function
that represents an edge-on disc with multiple radial truncations.
Inner parts of the galaxies were excluded from the fits to avoid
the CMC. Here, we are interested in both CMCs and discs, so
we extend our fits all the way to the centre. The fitted function
is the superposition of a function describing the disc and one
describing the CMC:
J(R) = Jdisc(R) + JCMC(R), (1)
where R is the projected distance from the centre of the galaxy
in the direction of the disc plane.
We assume the disc face-on luminosity profile Idisc to be ei-
ther exponential or that of a truncated disc (Erwin et al. 2008,
CO12) depending on the galaxy. The face-on luminosity profile
of a disc with multiple truncations is
I(r′) = S I0 e−r′/h1
i=n∏
i=2

[
1 + eαi−1,i
(
r′−r′i−1,i
)] 1
αi−1,i
(
1
hi−1 −
1
hi
) , (2)
where r′ is the 2D radius in the plane of the galaxy, I0 is the
central surface brightness, S is a normalization factor, n is the
number of exponential sections, hi is the scale length of the ex-
ponential section i, r′i−1,i is the break radius between the sections
1 http://irsa.ipac.caltech.edu/
2
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
Table 1. Properties of the galaxies in the sample
ID vc T d M3.6µm(AB) CMC ID vc T d M3.6µm(AB) CMC
(km/s) (Mpc) (km/s) (Mpc)
(1) (2) (3) (4) (5) (6) (1) (2) (3) (4) (5) (6)
ESO 157-49 115 3.0 23.8 -19.23 U NGC 4081 103 1.0 30.2 -19.93 N
ESO 240-11 276 4.8 40.8 -21.94 C NGC 4111 180 -1.4 15.6 -20.70 C?
ESO 292-14 131 6.5 28.7 -19.21 N NGC 4330 126 6.3 20.4 -19.48 N
ESO 346-1 127 5.1 28.2 -18.87 P NGC 4359 110 5.0 16.6 -18.29 N
ESO 440-27 132 6.6 20.9 -19.70 P NGC 4437 149 6.0 11.1 -20.32 P
ESO 443-21 164 5.8 47.5 -20.75 U NGC 4565 250 3.2 11.8 -21.83 CX?
ESO 466-14 124 3.8 47.7 -19.48 N NGC 4607 86 3.4 22.2 -19.78 U
ESO 469-15 104 3.3 27.4 -18.63 N NGC 4747 102 7.2 12.3 -18.35 N
ESO 533-4 159 5.2 40.2 -20.52 U NGC 5470 121 3.1 24.7 -19.55 N
ESO 544-27 125 3.4 38.6 -19.27 U NGC 5529 301 5.3 44.1 -22.20 UX
IC 217 104 5.9 26.0 -19.02 N NGC 5981 225 4.3 29.2 -20.48 U
IC 610 152 3.9 38.1 -20.65 U NGC 6010 158 0.4 21.6 -20.11 CU?
IC 1197 100 6.0 25.3 -18.73 N NGC 7347 127 4.6 38.7 -19.96 U
IC 1553 128 5.2 33.4 -19.65 U PGC 13646 179 5.1 34.4 -20.80 U
IC 1711 193 3.0 43.3 -21.19 CX? PGC 28308 180 6.8 75.1 -21.72 U
IC 1913 79 3.4 21.0 -17.43 N PGC 30591 122 6.8 60.3 -20.27 N
IC 2058 109 6.5 21.3 -18.48 N PGC 32548 103 -0.2 49.6 -18.59 CU
IC 2135 113 6.0 24.6 -20.08 N PGC 52809 112 5.9 24.2 -19.89 P
IC 5176 180 4.5 26.4 -20.86 U UGC 903 182 3.9 38.6 -21.23 U
NGC 489 191 2.7 43.3 -20.93 C? UGC 1970 109 5.9 36.0 -19.08 N
NGC 522 190 4.1 40.0 -21.03 UX UGC 5347 104 6.5 50.8 -19.03 U
NGC 678 196 3.0 27.9 -21.06 CX? UGC 5689 127 6.4 58.3 -20.01 U
NGC 1032 261 0.4 35.6 -21.76 C? UGC 5958 83 4.0 29.4 -18.20 N
NGC 1163 147 4.1 33.6 -19.84 U UGC 6526 90 7.0 32.3 -18.97 N
NGC 1422 72 2.3 21.5 -18.48 N UGC 7086 136 3.1 37.0 -20.22 U
NGC 1495 114 5.0 16.7 -18.79 N UGC 8737 169 4.0 39.1 -20.83 U
NGC 1596 194 -2.0 15.6 -20.22 C? UGC 9448 123 3.2 46.7 -19.45 U
NGC 2732 162 -2.0 27.5 -20.52 P? UGC 9665 146 4.0 39.6 -20.32 U
NGC 3098 173 -1.5 20.3 -19.72 CX? UGC 10043 160 4.1 48.9 -20.75 C?
NGC 3279 174 6.5 36.7 -20.97 N UGC 10288 180 5.3 30.8 -20.51 U
NGC 3454 103 5.5 26.6 -19.26 U UGC 10297 112 5.1 38.8 -18.89 N
NGC 3501 146 5.9 23.8 -20.11 U UGC 12518 176 3.0 33.9 -19.65 U
NGC 3592 94 5.3 26.6 -18.80 U UGC 12692 113 3.9 49.6 -19.73 N
NGC 3600 111 1.0 17.0 -18.29 C? UGC 12857 117 4.0 35.4 -19.66 U
NGC 3628 226 3.1 11.3 -21.73 PX
Notes. ID (Col. 1) refers to the galaxy name and vc (Col. 2) refers to the galaxy circular velocity from the EDD, except for ESO 466-14
(Mathewson & Ford 1996), NGC 1596 (Williams et al. 2010), NGC 2732 (Simien & Prugniel 1997), NGC 3098 (Balkowski & Chamaraux 1981),
and NGC 4111 (Simien & Prugniel 1998). T (Col. 3) indicates the galaxy stage from HyperLeda, d (Col. 4) its distance from NED, and M3.6µm(AB)
(Col. 5) its absolute magnitude in 3.6µm from the S4G Pipeline 3 fluxes and Col. 4. CMC (Col. 6) indicates the kind of CMC as classified by the
authors of this paper (see Section 5). “C” stands for classical bulge, “P” for pseudo-bulge, “U” for unresolved CMC, “X” for boxy or X-shaped
inner isophotes, and “N” for none. Doubtful cases are indicated by a question mark “?”. Light from boxy or X-shaped structures has been treated
as belonging to the disc in our analysis.
i−1 and i, and αi−1,i controls the sharpness of the break. As done
in CO12 and Laine et al. (2014), we fix αi−1,i = 0.′′5 for all the
breaks. This value is typical of the sharpness of the breaks in
Erwin et al. (2008). The maximum number of fitted exponential
sections is nmax = 4. We integrate the face-on luminosity profiles
along the line of sight to obtain edge-on luminosity profiles as
follows:
Jdisc(R) = Fdisc
∫ s=sf
s=−sf Idisc
(√
R2 + s2
)
ds
∫ R=5Rf
R=−5Rf
∫ s=sf
s=−sf Idisc
(√
R2 + s2
)
ds dR
, (3)
where the parameter Fdisc describes the total flux from the disc.
To ease the integration, we consider the discs to be truncated at
a large 2D radius. We select this radius to be r′f = 5Rf where
Rf is the projected radius over which the disc fits were produced
in CO12, s is the coordinate along the line of sight, and sf =√
(5Rf)2 − R2. The free parameters of the fit are Fdisc, the scale
lengths of the exponential disc sections, and the break radii.
We assume that the 2D surface brightness of the CMC fol-
lows a Se´rsic profile (Se´rsic 1963)
ICMC(r) = ICMC(re)e[−bn(r/re)1/n−1] (4)
where r is the 2D radius in the plane of the sky, re is the half-
light radius, and n is the Se´rsic index that describes the shape
of the profile. We use bn ≈ 1.999n − 0.327 (Capaccioli 1989;
Graham & Colless 1997; Graham & Driver 2005). In all cases
re ≪ Rf and zu < Rf , so we consider the CMCs to be truncated
at a radius r = Rf to ease integration. Hence,
JCMC(R) = FCMC
∫ z=min(zu,zf)
z=−min(zu,zf) ICMC
(√
R2 + z2
)
dz
∫ R=Rf
R=−Rf
∫ z=zf
z=−zf ICMC
(√
R2 + z2
)
dz dR
(5)
where z is the projected distance from the centre of the galaxy
in the direction perpendicular to the galaxy mid-plane, zf =
3
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
√
R2f − R2, and FCMC is the total flux from the CMC. The free
parameters of this function are FCMC, re, and n.
Because of the large number of free parameters involved, we
do not fit them all at the same time. Instead, we divide the fitting
process into two steps. In step one, we take the radial luminosity
profile and subtract the light corresponding to the disc as fitted in
CO12. In some cases, these disc fits have to be slightly changed
to ease the convergence (e.g., by adding inner disc sections in
regions that were not fitted in the previous paper due to the pres-
ence of the CMC). The residual is fitted with Equation 5 using
IDL’s curvefit. Then, in step two, we fit the full radial profile
with the function in Equation 1 using the JCMC(R) found in step
one and keeping it fixed. Then we repeat step one, but this time
we subtract the disc function obtained in step two and so on. As
in CO12, we used a Gaussian kernel with a 2.′′2 full width at half
maximum (FWHM) to take into account point spread function
(PSF) effects. We iterate until none of the free parameters varies
by more than 1%. In five cases with unresolved or barely re-
solved CMCs (NGC 3454, NGC 5981, UGC 7086, UGC 8737,
and UGC 9448), we fix n = 1 in order to prevent the Se´rsic index
from diverging.
An example of a radial luminosity profile and its fit are pre-
sented in Figure 1. The fits for all the galaxies with a detectable
CMC can be found in the Appendix A. The fits of radial lumi-
nosity profiles of galaxies with no detectable CMC are the same
as presented in the on-line material in CO12.
The typical formal errors in the fluxes of CMCs and discs
are ∼ 5%. An additional source of error is that our fitting func-
tion considers CMCs to be spherical. To check what would
happen if CMCs were in fact disky or boxy, we first create
bulges with non-circular isophotes using the generalized el-
lipses defined in Athanassoula et al. (1990). We then build a
grid of oblate bulges with axis ratios b/a = [0.6, 0.8, 1.0],
diskyness/boxyness parameter c = [1.7, 2, 2.3] (c is defined in
Athanassoula et al. 1990), Se´rsic index n = [1, 2, 3, 4], effective
radii re = [0.9, 1.8, 3.6, 7.3] (in units of the PSF FWHM), and
brightnesses LCMC = [0.05, 0.1, 0.2, 0.4] (in units of Ldisc). Then
we embed these bulges into single-exponential discs with scale
lengths h = [6.8, 13.6, 20.5] (in the same units as re). The total
number of models in the grid is 1728. Then we fit the radial lu-
minosity profiles of our synthetic galaxies with Equation 1. We
find that the fitted CMC flux is not very sensitive to the oblate-
ness and boxiness of the CMC (typical flux errors are smaller
than 10%). However, we find that when the CMC has an effec-
tive radius comparable to the scale length of the disc, the CMC
component often fits part of the disc, which causes the CMC flux
to be overestimated. When this happens, the CMC effective ra-
dius is also overestimated by a large amount (a factor typically
much larger than five). These cases are rare, because even in the
case of S0 galaxies, the median ratio between the bulge effective
radius and the disc scale length is re/h ∼ 0.2 (Laurikainen et al.
2010). Also, we are confident that we would be able to detect
these largely overestimated effective radii in the real galaxy fits.
Once these cases are not considered, the typical uncertainty in
the fitted CMC fluxes (accounting for the formal errors) is of the
order of 10%. The typical error in the fitted disc fluxes is 5%.
The fits provide estimates of the fractions of the flux emitted
by the CMC and the disc. Also, from CO12, we know the frac-
tion of disc light that corresponds to the thin and thick disc com-
ponents. By making assumptions on the value of mass-to-light
ratios of the discs and the CMC, we can estimate the masses of
each of these components.
Here, as in CO12, we assume that the ratio of 3.6µm mass-
to-light ratios of the thick and the thin discs is ΥT/Υt = 1.2.
We assume that Υt = 0.55, so the global mass-to-light ra-
tio of the disc is compatible with the Υ ≈ 0.6 value from
Meidt et al. (2014). There is some evidence that the Milky Way
CMC might have a similar chemical history to that of the thick
disc (Mele´ndez et al. 2008). Therefore, we adopt a CMC mass-
to-light ratio ΥCMC = ΥT. This also agrees with the hypothesis
(discussed in Section 5) that CMCs and thick discs formed at
the same time. We assume that M⊙,3.6µm(AB) = 6.06 (Oh et al.
2008).
As in CO12 we obtain atomic gas masses from the
HyperLeda fluxes with the method proposed by Zwaan et al.
(1997) and by multiplying the atomic gas mass by a 1.4 factor
to account for He and metals. Fluxes are not available for two
galaxies, namely ESO 466-14 and PGC 30591. In what follows,
we considered Mg = 0 for these two galaxies. Molecular gas is
not considered.
The masses of the gas disc, thin disc, the thick disc, and the
CMC are designated as Mg, Mt, MT, and MCMC, respectively.
4. Results
4.1. Component masses as a function of vc
In Figure 2 we plot the gas disc, thin disc, the thick disc, and the
CMC masses as a function of the galaxy circular velocity. The
galaxy circular velocity is a good proxy for the total galaxy mass.
In most cases, we use circular velocities based on gas rotation
from the Extragalactic Distance Database (EDD; Courtois et al.
2009). In two cases for which EDD velocities are not available
(ESO 466-14, NGC 3098), we use data from Mathewson & Ford
(1996) and Balkowski & Chamaraux (1981). However, in the
case of perturbed or truncated gas discs, their circular veloci-
ties might not be representative of the true mass of the system.
We compared HI circular velocities with those from stellar ve-
locities based on absorption lines in Simien & Prugniel (1997,
1998) (four galaxy overlap) and those from dynamical models
in Williams et al. (2010) (four galaxy overlap). We find them to
be the same within 10% except for NGC 1596, NGC 2732, and
NGC 4111. These three galaxies are S0s which are likely not to
have much gas. For these three galaxies, we adopt the circular
velocities not based on gas.
In Figure 2, the error bars for the CMC masses correspond
to a 10% fitting uncertainty. The thin and thick disc masses are
calculated using the fraction of light corresponding to the disc
and by using the ratio of thick to thin disc masses from CO12
(MT/Mt). The errors in the thin and thick disc masses come
from a quadratic sum of two factors. The first one is the 5% fit-
ting uncertainty in the fraction of light assigned to the combined
disc from our luminosity profile fit. The second one is a fixed
20% of the thin or thick disc mass. This factor is a reasonable
estimate of the error caused by the biases in the thin/thick disc
fits (Section 4.3.2 in CO12). The fitted Mt, MT, and MCMC val-
ues are presented in Table 2. This table also includes the atomic
gas mass estimates,Mg. The errors in the gas masses come from
the flux error estimates from HyperLeda.
At least two kinds of CMCs may have their measured masses
different from the real value for reasons other than those ac-
counted by the error bars. First, some CMCs might actually be
nuclear clusters with a small ΥCMC. This is most likely to happen
in unresolved CMCs which are marked in Figure 2 with green
circles. Here we have marked as unresolved those CMCs with
re < 3′′. Second, in seven cases, the CMC is superposed to a
4
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
Table 2. Masses of the galaxy components
ID Mg Mt MT MCMC ID Mg Mt MT MCMC
(107 M⊙) (107 M⊙) (107 M⊙) (107 M⊙) (107 M⊙) (107 M⊙) (107 M⊙) (107 M⊙)
ESO 157-49 190 ± 10 350 ± 70 380 ± 80 54 ± 5 NGC 4081 172 ± 6 800 ± 200 600 ± 100 −
ESO 240-11 5300 ± 300 7000 ± 1000 1200 ± 300 1400 ± 100 NGC 4111 200 ± 20 1500 ± 300 500 ± 100 1000 ± 100
ESO 292-14 650 ± 40 440 ± 90 320 ± 60 − NGC 4330 110 ± 10 480 ± 100 500 ± 100 −
ESO 346-1 160 ± 20 250 ± 50 190 ± 40 130 ± 10 NGC 4359 138 ± 9 180 ± 40 150 ± 30 −
ESO 440-27 1170 ± 90 390 ± 80 380 ± 80 480 ± 50 NGC 4437 730 ± 40 1500 ± 300 430 ± 90 110 ± 10
ESO 443-21 810 ± 50 2200 ± 400 700 ± 100 140 ± 10 NGC 4565 2100 ± 200 6000 ± 1000 1700 ± 300 470 ± 50
ESO 466-14 − 600 ± 100 390 ± 80 − NGC 4607 118 ± 7 700 ± 100 600 ± 100 23 ± 2
ESO 469-15 500 ± 50 240 ± 50 210 ± 40 − NGC 4747 106 ± 5 160 ± 30 200 ± 40 −
ESO 533-4 680 ± 20 1500 ± 300 1000 ± 200 19 ± 2 NGC 5470 320 ± 20 600 ± 100 400 ± 80 −
ESO 544-27 420 ± 30 480 ± 100 280 ± 60 34 ± 3 NGC 5529 3800 ± 100 7000 ± 1000 4100 ± 800 1000 ± 100
IC 217 300 ± 10 380 ± 80 250 ± 50 − NGC 5981 520 ± 50 1700 ± 300 500 ± 100 140 ± 10
IC 610 180 ± 20 1800 ± 400 900 ± 200 93 ± 9 NGC 6010 27 ± 1 1100 ± 200 350 ± 70 250 ± 20
IC 1197 320 ± 10 160 ± 30 350 ± 70 − NGC 7347 1200 ± 200 900 ± 200 500 ± 100 45 ± 4
IC 1553 250 ± 20 800 ± 200 220 ± 50 58 ± 6 PGC 13646 660 ± 20 2300 ± 500 700 ± 100 220 ± 20
IC 1711 520 ± 20 1800 ± 400 1200 ± 200 1800 ± 200 PGC 28308 1400 ± 100 4800 ± 1000 1900 ± 400 910 ± 90
IC 1913 165 ± 5 90 ± 20 60 ± 10 − PGC 30591 − 1400 ± 300 600 ± 100 −
IC 2058 430 ± 20 160 ± 30 250 ± 50 − PGC 32548 690 ± 20 230 ± 50 120 ± 20 83 ± 8
IC 2135 480 ± 30 700 ± 100 1000 ± 200 − PGC 52809 1260 ± 60 600 ± 100 700 ± 100 250 ± 30
IC 5176 2300 ± 200 2200 ± 400 1100 ± 200 84 ± 8 UGC 903 1340 ± 50 3300 ± 700 1300 ± 300 160 ± 20
NGC 489 246 ± 9 2200 ± 500 700 ± 100 760 ± 80 UGC 1970 450 ± 10 440 ± 90 220 ± 40 −
NGC 522 110 ± 7 2500 ± 500 1400 ± 300 140 ± 10 UGC 5347 1000 ± 100 420 ± 80 180 ± 40 31 ± 3
NGC 678 500 ± 30 1800 ± 400 700 ± 100 1800 ± 200 UGC 5689 880 ± 30 1000 ± 200 500 ± 100 38 ± 4
NGC 1032 26 ± 1 3700 ± 800 1500 ± 300 2900 ± 300 UGC 5958 110 ± 10 150 ± 30 160 ± 30 −
NGC 1163 560 ± 30 700 ± 100 480 ± 100 170 ± 20 UGC 6526 120 ± 10 350 ± 70 260 ± 50 −
NGC 1422 79 ± 3 180 ± 40 220 ± 40 − UGC 7086 900 ± 40 1100 ± 200 700 ± 100 67 ± 7
NGC 1495 160 ± 10 330 ± 70 180 ± 40 − UGC 8737 174 ± 6 1900 ± 400 1300 ± 300 94 ± 9
NGC 1596 210 ± 20 700 ± 100 280 ± 60 1100 ± 100 UGC 9448 820 ± 40 600 ± 100 370 ± 70 9 ± 1
NGC 2732 35 ± 5 1000 ± 200 700 ± 200 810 ± 80 UGC 9665 1700 ± 200 1500 ± 300 460 ± 90 120 ± 10
NGC 3098 140 ± 20 500 ± 100 260 ± 50 510 ± 50 UGC 10043 2600 ± 200 1500 ± 300 1100 ± 200 660 ± 70
NGC 3279 360 ± 30 2400 ± 500 1400 ± 300 − UGC 10288 1800 ± 60 1400 ± 300 900 ± 200 180 ± 20
NGC 3454 150 ± 10 500 ± 100 260 ± 50 21 ± 2 UGC 10297 410 ± 20 360 ± 70 200 ± 40 −
NGC 3501 520 ± 20 1000 ± 200 700 ± 100 38 ± 4 UGC 12518 260 ± 30 500 ± 100 360 ± 70 300 ± 30
NGC 3592 82 ± 8 240 ± 50 270 ± 50 19 ± 2 UGC 12692 290 ± 20 800 ± 200 400 ± 80 −
NGC 3600 680 ± 50 100 ± 20 130 ± 30 120 ± 10 UGC 12857 980 ± 80 800 ± 200 290 ± 60 37 ± 4
NGC 3628 740 ± 30 3000 ± 600 4400 ± 900 650 ± 70
Notes. Mg from HyperLeda. The remaining masses are from fits in this paper and CO12.
boxy or X-shaped structure which could be caused by an edge-
on bar. These structures are detected visually. Although the light
of these structures seems to be mostly fitted by the inner disc
sections, they could slightly perturb the fitting of the CMC. In
Figure 2, we indicate CMC fits that could be perturbed by X-
shaped structures with crosses. Because these structures are fit-
ted as part of the disc, their light is assigned into the thin and the
thick disc components in the same proportions as for the rest of
the disc light.
CMCs that are resolved are visually classified as classical
bulges or pseudobulges according to their morphology. CMCs
protruding the disc or with round isophotes (axis ratio b/a & 0.6)
are considered to be classical bulges. This is because classi-
cal bulges are generally thought to be pressure supported, while
pseudobulges are often thought to have a significant amount of
rotation and therefore be more flattened than classical bulges.
All protruding CMCs have round isophotes. This classification
is uncertain in many cases due to its subjectivity, and also to the
fact that some of the CMCs are poorly resolved.
The CMC classifications are presented in Table 1 and doubt-
ful cases are indicated with a question mark.
The masses of the thin discs, the thick discs, and the CMCs
increase with vc. Thick discs generally have masses in the range
109M⊙ − 1010M⊙ for low-mass galaxies (vc . 120 km s−1)
but they can be as massive as several times 1010M⊙ for galax-
ies with a larger mass. As shown in Comero´n et al. (2011a) and
CO12, thin discs have masses similar to those of thick discs for
vc < 120 km s−1 and are on average more massive than thick
discs for galaxies with larger rotation velocities. CMCs are al-
most always less massive than both the thin and the thick discs.
All but five of the 23 galaxies without a CMC are found in galax-
ies with vc . 120 km s−1. Gas disc masses have little or no de-
pendence with vc and are most often in the 109M⊙ − 1010M⊙
range.
4.2. Mass ratios of the different components
In this paper, we consider that the dynamically “hot compo-
nent” is the sum of the CMC and the thick disc. This is based
on the observation that young galaxies are highly turbulent
(Fo¨rster Schreiber et al. 2009; Kassin et al. 2012), which means
that their gas discs are thicker than today’s gas discs. Stars that
formed in this turbulent gas should have shared the same large
thickness as the gas, just like today’s young stars have the same
thickness as today’s gas. This means that all stars in a young
galaxy were in a relatively thick disc; there was no thin disc
component for either the stars or the gas (Bournaud et al. 2009).
Consequently, any angular momentum exchange inside the disc
5
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
50 100 150 200 250 300
vc (km s
-1)
107
108
109
1010
1011
M
 (
M
O •
 )
Gas disc
Thin disc
Thick disc
CMC
Hot components
Cold components
Unresolved CMC
Boxy inner structure
Classical bulge
Pseudobulge
No detectable CMC
Fig. 2. Mass of the gas discs (deep blue circles), thin discs (blue circles), thick discs (orange circles), and CMCs (red circles)
as a function of the galaxy circular velocity. Red arrows indicate galaxies for which no significant CMC is found. Red circles in
large green circles indicate unresolved CMCs, red circles in large dark orange circles indicate classical bulges, and simple red
circles correspond to pseudobulges. Deep blue arrows indicate galaxies for which the 21-cm flux is not available. Crosses indicate
boxy/X-shaped inner structures that are superimposed onto a given CMC. The vertical dashed line corresponds to vc = 120 km s−1.
at that time, leading to disc contraction or clump migration to the
centre, will have made a thick central concentration. This differs
from disc migration in a later stage cool disc, which can make a
thin inner concentration. It also differs from the later formation
of a bar out of a cool inner stellar disc, and from the thickening
produced by the bar through vertical instabilities and bar buck-
ling. Thus we consider both thick discs and CMCs formed in an
early disc with a large scale height compared to that of the thin
disc. The dynamically “cold component” is made of the thin disc
and the gas disc (for which the atomic gas mass is a lower mass
boundary).
Even though some of our CMCs might actually be unre-
solved nuclear clusters unrelated to classical bulges and pseu-
dobulges, we are not making a large error at measuring the mass
of the hot component. This is because the unresolved CMCs
usually are the least massive ones. Also, we are ignoring the
fact that at least part of the mass of pseudobulges is a prod-
uct of secular evolution and are therefore a disc-like structure.
We visually checked whether the six pseudobulges identified in
the sample were thicker than the thin disc and could therefore
be considered a “hot component”. We find that three of them
(NGC 3628, NGC 4437, and PGC 52809) have thicknesses com-
parable to that of the thin disc. The three have relatively small
masses (MCMC/(MCMC +MT +Mt +Mg) < 0.1). These three
“flat” CMCs might be the product of secular evolution and actu-
ally be “disky” pseudobulges.
Previous studies showed that MT/Mt decreases with vc
(Yoachim & Dalcanton 2006; Comero´n et al. 2011a, CO12). We
show this trend in the top panel of Figure 3. A similar but milder
trend is found for the ratio between the thick disc mass and the
cold component mass,MT/(Mt+Mg) (second panel in Figure 3,
also shown in Yoachim & Dalcanton 2006, and in CO12). These
two trends might be related to that found in a sample of mod-
erately inclined galaxies at intermediate redshifts by Sheth et al.
(2012): their Figure 4 relates kinematic properties of the galax-
ies (from optical spectroscopy) to their masses and shows that
low-mass galaxies tend to be dynamically hotter than those with
a larger mass. In other words, the smaller the mass, the larger
the ratio σ/vc, where σ is the velocity dispersion. The trend pre-
sented here and the one displayed in Sheth et al. (2012) might be
related because the component of the velocity dispersion perpen-
dicular to the galaxy plane (σz) and the scale height (z0) have a
positive correlation. Sheth et al. (2012), however, measured gas
kinematics and not stellar kinematics.
The third panel of Figure 3 shows the ratio of the masses of
the CMC and the thick disc as a function of vc. It shows a trend
between MCMC/MT and vc. This indicates that the hot compo-
nent is more centrally concentrated in more massive galaxies.
In the fourth and fifth panels of Figure 3, we show that the
ratio between the mass of the hot component (sum of the thick
disc and the CMC) and that of the thin disc and the ratio between
the mass of the hot and the cold components have no clear trend
with the circular velocity, although the plots have a large scatter.
The lack of a clear trend in these ratios might be caused by large
intrinsic scatter and fitting uncertainties which might mask any
correlation. However, a more interesting possibility is that these
ratios are indeed mass-invariant. This is because, as seen in the
first, the second and the third panels, MT/Mt and MT/(Mt +
Mg) decrease with vc whereas MCMC/Mt increases with vc in a
way that the two effects roughly cancel each other.
6
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
Unresolved CMC
Boxy inner structure
Classical bulge
Pseudobulge
No detectable CMC
      
      
0.1
1.0
M
T
/M
t
 
 
      
      
0.1
1.0
M
T
/(
M
t+
M
g
)
 
 
      
      
0.1
1.0
(M
T
+
M
C
M
C
)/
M
t
 
 
50 100 150 200 250 300
vc (km s
-1)
      
      
0.1
1.0
(M
T
+
M
C
M
C
)/
(M
t+
M
g
)
 
 
10-3
10-2
10-1
100
M
C
M
C
/M
T
      
      
Unresolved CMC
Boxy inner structure
Classical bulge
Pseudobulge
No detectable CMC
Fig. 3. Mass of the thick disc divided by that of the thin disc
(first panel), mass of the thick disc divided by that of the cold
component (second panel), mass of the CMC divided by that of
the thick disc (third panel), and mass of the hot component di-
vided by that of the thin disc (fourth panel), and mass of the hot
component divided by that of the cold component (fifth panel),
as a function of the circular velocity. The symbols indicate the
properties of the CMC of each galaxy and follow the same con-
ventions as in Figure 2 except for the triangles in the first, sec-
ond, third, and fourth and the bottom panels that indicate galax-
ies with no detectable CMC. In the second and fifth panels, the
two galaxies with no available 21-cm flux measurement are con-
sidered to have Mg = 0.
5. Discussion
5.1. Star formation rates
The star formation intensity (SFI) in a galaxy disc (the star
formation rate, SFR, per unit of area) can be expressed as
˙Σ∗ = ǫorbΣ/torb, where Σ is the total surface density, ǫorb is
the fraction of mass converted into stars in a galactic rota-
tion, and torb is the rotation period of the studied region (here
we use the formalism from Krumholz et al. 2012). In a single-
component disc in equilibrium (Spitzer 1942) the scale height
is z0 = σ2z/(πGΣ). In a gas-rich disc, the Toomre parame-
ter is Q = √2(β + 1)σΩ/(πGΣ), where β describes the slope
of the rotation curve, and Ω is the angular rotation velocity
(Krumholz et al. 2012). If we assume that σz is porportional to
σ, in a gas-rich galaxy with a single-component disc
˙Σ∗ ∝ ǫorbtorb
2(β + 1)Ω2
πGQ2 z0 (6)
and the SFI is proportional to the scale height ( ˙Σ∗ ∝ z0).
Dynamically hot components, if born already hot, must have
formed at a time when the SFI was larger than it is today. Thick
discs being a consequence of turbulence caused by a high SFI
was already proposed by Lehnert et al. (2014). High turbulence
could also be caused by accretion energy (Elmegreen & Burkert
2010) and by gravitational instabilities and stirring from the
massive clumps that form in those instabilities, even without star
formation feedback (Bournaud et al. 2009).
The conditions for a large scale height were fulfilled during
the early disc phase of galaxies when the cosmic star formation
rate was larger than it is now (see Madau & Dickinson 2014,
for a recent review). This epoch overlaps with that of clumpy
discs described in Elmegreen & Elmegreen (2006). The size of
a clump formed by gravitational instabilities is about the Jeans’
length, and this is about the same as the disc thickness because
both depend on the ratio of the square of the velocity disper-
sion, σ, to the disc mass column density, Σ, so the inverse Jeans’
wavenumber is λJ = σ2/ (πGΣ). Thus the observed large sizes of
star-forming clumps in young galaxies serve as an indirect mea-
sure for large disc thicknesses: clumpy discs should also be thick
discs (Elmegreen et al. 2009). If the hot components formed at
that time, then the ratio the hot component mass (MT +MCMC)
to the thin disc mass (Mt) should be equal to the ratio of the SFR
multiplied by the duration of star formation for the hot and the
cool phases:
MT +MCMC
Mt
≈
˙Σ∗h × th
˙Σ∗t × tt
. (7)
From Figure 3 we find that (MT +MCMC)/Mt ∼ 1. Therefore,
because ˙Σ∗h > ˙Σ∗t (this is a consequence of ˙Σ∗ ∝ z0), the time
taken to build the hot component was shorter than the time to
build the cool thin disc, th < tt. Assuming that th ∼ 1 Gyr (the
approximate duration of the disc clumpy phase; Bournaud et al.
2007) and tt ∼ 10 Gyr (the time elapsed since then), we find
˙Σ∗h/ ˙Σ∗t to be ∼ 10 with a large spread. However, this ratio
is highly dependent on the duration of the intense star forma-
tion phase. Longer or shorter hot component formation times
can change it significantly. To exemplify this, we can calculate
this ratio for the Milky Way by using star formation histories
in the literature. Snaith et al. (2014) claim that for the Galaxy
MT/Mt ∼ 1. They find that the thick disc formed in an intense
star formation phase that lasted for ∼ 3.5 Gyr and that the thin
disc has formed later during ∼ 9 Gyr. Assuming that the mass of
7
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
the CMC of the Milky Way is small compared to that of the disc
components, we obtain ˙Σ∗h/ ˙Σ∗t ∼ 2.5.
5.2. Thick disc and CMC formation
The observations presented here suggest that about half of the
stars in a galaxy are in a hot and thick component. This is a
sensible fraction considering the variation of the star forma-
tion rate and mass build-up over a Hubble time. According
to the toy model in Dekel et al. (2013), half of the mass in a
galaxy is accreted before z ∼ 0.9, when the accretion rate ex-
ceeded by 2.4 times the current rate (M ∝ exp(−0.79z) and
˙M ∝ exp(−0.79z)(1+z)5/2; the universe is assumed to be matter-
dominated for most of a galaxy’s accretion history). In an equi-
librium situation, and assuming that the surface of the galaxy
does not change dramatically, the SFI is roughly proportional
to the accretion rate (Dekel et al. 2013). With the assumption
of a constant Toomre Q and pressure equilibrium in the disc
(Krumholz et al. 2012), a SFI that scales with the gas column
density also scales with the velocity dispersion and the scale
height. Therefore, one would expect the scale height for instanta-
neous star formation to scale with the accretion rate. This means
that by the time the first half of the stellar disc formed, at z ∼ 0.9,
the thickness had always been larger than 2.4 times the present
thickness, just from the higher accretion and SFI at these ear-
lier times. The second half of the disc formed thinner because
of the lower SFI. If we identify the first half of the assembly
with the thick disc and CMC components of today’s galaxies
and the second half with the thin disc and gas components, then
this minimum theoretical thickness for the first half is consis-
tent with our observations in Comero´n et al. (2011a), where the
thick-to-thin scale height ratio is zT/zt ∼ 4, increasing slightly
for earlier types. In this assembly model, the old and thick discs
of spiral galaxies are direct results of the high accretion rates at
earlier cosmic times.
This picture is consistent with the formation of the thick
disc and CMC in a relatively short time in a clumpy disc
(Bournaud et al. 2007, 2009) at high redshift and a subsequent
slow formation of the thin disc. It is also consistent with the
α-enhancement found for the Galaxy CMC and thick disc
(Mele´ndez et al. 2008), and with the age differences between ex-
tragalactic thin and thick discs found by Yoachim & Dalcanton
(2008). Different simulations have shown that the inward mi-
gration of clumps in clumpy discs can build classical bulges
(Elmegreen et al. 2008; Perez et al. 2013) and/or pseudobulges
(Inoue & Saitoh 2012, 2014). The mechanisms causing the loss
of angular momentum of clumps (ultimately leading to their
merger and the formation of the CMC) would then be more effi-
cient in the most massive galaxies, where the CMC mass is the
largest.
The exact relationship between the SFI and the SFR is hard
to asses. Lehnert et al. (2014), in their Milky Way SFI study,
assume that for the Milky Way the thick disc scale length is
shorter than the thin disc scale length as advocated by many
(Bensby et al. 2011; Bovy et al. 2012; Cheng et al. 2012) using
chemical properties to divide stellar populations. Thus, the re-
duction in the SFI would be both a consequence of the increase
in the surface of the disc and the reduction of the gas accretion.
Interestingly, studies based on structural decompositions show
the opposite, namely that the scale length of the thick disc is
larger than that of the thin disc (Juric´ et al. 2008; Chang et al.
2011; Polido et al. 2013; Lo´pez-Corredoira & Molgo´ 2014).
Regarding other galaxies, structural decompositions show that
in general thick discs have large scale lengths compared to thin
discs (Yoachim & Dalcanton 2006, CO12), which is in uncom-
fortable tension with some of the studies regarding the Milky
Way morphology. In this case, the SFI evolution would be driven
by two processes working in oposite directions, namely a reduc-
tion in the accretion rate and a reduction of the radius within
which the star formation is found. The fact that thin discs are
indeed thin shows that the effect of the decreased gas accretion
would be more important than that of the shrinking disc.
The mean CMC-to-total mass ratio in our study of edge-on
galaxies is B/T = 0.09. This is comparable to the observed
B/T in not very inclined Sab-Sc galaxies (Laurikainen et al.
2010), but is smaller than that found in the simulation papers
where a CMC is formed in a clumpy disc. Elmegreen et al.
(2008) fit a double-disc to the outer parts of galaxies and define
the CMC as the central excess of light. They find B/T values
in the range between 0.12 and 0.36. Perez et al. (2013) make
Se´rsic+exponential disc radial fits to their models and obtain
B/T values in the range between 0.24 and 0.39. Inoue & Saitoh
(2014) also use Se´rsic+exponential fits and obtain B/T in the
range between 0.23 and 0.35. A reason for the difference be-
tween our B/T values and those in simulations is that the
latter correspond to the ratio before the thin disc formed.
Hence, it corresponds to MCMC/(MCMC +MT) rather than to
MCMC/(MCMC+MT+Mt). The averageMCMC/(MCMC+MT)
vtalue for our sample is 0.19, which is closer to what is found in
simulations.
5.3. Are all CMCs hot components?
A caveat of our study is that we consider the whole mass of
CMCs to be part of a hot component made at high redshift. This
assumption may not always be true, especially for pseudobulges,
often described as the consequence of a secular process.
A recent study of the colours of pseudobulges (defined
to have n < 2.5) in a sample of isolated galaxies has
shown that ∼ 2/3 of them are found in the red sequence
(Ferna´ndez Lorenzo et al. 2014). They show that those red
pseudobulges have colours compatible with a single star for-
mation burst 8 Gyr ago (z ∼ 1 in a standard cosmology).
MacArthur et al. (2009) also find for a sample of CMCs with
n < 2.5 that the fraction of the stellar mass in the CMC with ages
> 10 Gyr is at least 50%. However, pseudobulges have a frac-
tion of mass built secularly either in a continuous or an episodic
way: Fisher et al. (2009) define the CMC growth time as its stel-
lar divided by the present-day SFR and find that ∼ 50% of their
pseudobulges have a growth time of less than 10 Gyr (the cur-
rent SFR is larger than the past average SFR which is the CMC
mass divided by a Hubble-Lemaıˆtre time) and thus still have a
significant SFR. In their study, pseudobulges are defined to have
n < 2 and/or disc morphology features such as spiral arms.
Another piece of evidence for the building of a large frac-
tion of the CMC mass during the large SFI phase is found in
van Dokkum et al. (2013). They find that after z ∼ 1, the SFR
of Milky Way-like galaxies drops, and that this drop is stronger
in the central 2 kpc. Therefore, while the disc continues to be
slowly built after z ∼ 1, the growth of the CMC has almost
stopped.
Whatever is the fraction of mass in CMCs built at high red-
shift, the total mass of the CMC is only an upper limit to the
mass of the CMC built as a hot component at high z.
8
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
6. Conclusion
In this paper, we study vertically integrated radial luminosity
profiles from mid-infrared images of 69 edge-on galaxies from
the S4G. We produce photometric decompositions of the profiles
to distinguish the light that comes from the central mass concen-
tration (CMC) from that coming from the disc. Here we assume
CMCs to be classical bulges and pseudobulges, and we exclude
boxy and X-shaped features because we consider them to be re-
lated to disc features (bars). Hence, the light from boxy and X-
shaped features is assigned to discs in our analysis. We combine
this information with that in CO12 and assumptions on compo-
nents’ mass-to-light ratios to obtain the masses of the CMCs,
thin discs, and thick discs (MCMC, Mt, and MT, respectively).
We obtain atomic gas disc masses, Mg, from 21-cm fluxes in
HyperLeda. We consider the sum of the CMC and the thick disc
to be the dynamically hot component of a galaxy and the sum of
the gas disc and the thin disc to be the dynamically cold compo-
nent.
We confirm a previous result that the ratio of the thick to
thin disc masses, MT/Mt, decreases with increasing vc. A sim-
ilar but milder trend is found between MT/(Mt + Mg) and
vc. However, the ratio of hot component to thin disc masses,
(MCMC +MT)/Mt, and the ratio of the hot component to cold
component masses, (MCMC +MT)/(Mt +Mg), are roughly in-
dependent of the galaxy mass (although with a large scatter). We
also find that the larger vc, the more centrally concentrated the
hot component is (the ratio MCMC/MT becomes larger).
We suggest that our results are compatible with the two hot
components (the thick disc and the CMC) being built in a short
time during the rapid accretion phase of disc galaxies when the
discs were clumpy and the velocity dispersions, scale heights,
and star formation rates were all high. In this picture, the thin
disc slowly forms afterwards from gas accreted in cold flows.
Moreover, we suggest that the CMC formed by concentration of
the thick disc during this hot phase, and that is why both the thick
disc and the CMC together give a fixed fraction of the total disc
mass, and not just one or the other alone. We compare the ratio
of hot component to thin disc masses, (MCMC+MT)/Mt, in our
sample with the predictions from toy models. We estimate that
the SFI during the formation of the hot components was roughly
ten times larger than it was during the formation of the thin disc.
The value of the f SFI during the former phase relative to that of
the thin disc formation is a function of the duration of the high
star formation intensity epoch.
When combined with observations of galaxies at intermedi-
ate to high redshift, our results suggest that the hot disc phase
occurred during the high accretion phase of galaxy formation,
when the relative gas mass was high following the high accretion
rate, the star formation rate was high because of the large relative
gas mass, and the gas turbulent speed was high as a result of sev-
eral processes, including accretion energy input, violent gravita-
tional instabilities, stirring by massive clumps that form in these
instabilities, and feedback by rapid star formation. In this model,
the high turbulent speed for the gas is the single most important
reason for the thick disc and for any thick central mass concen-
tration that forms in it. Highly turbulent discs are automatically
thick because of their large pressures. Any stars that form in a
thick gas disc will share the high random speeds and thick distri-
butions of the gas. Bulge formation is a secondary step involving
angular momentum transfer in a disc that is already thick from
high turbulence.
Acknowledgements. We thank our referee, Matthew Lehnert, who helped to sub-
stantially improve the discussion section of this paper. We thank Simo´n Dı´az-
Garcı´a, Mirian Ferna´ndez Lorenzo, Martı´n Herrera-Endoqui, Joannah L. Hinz,
Luis C. Ho, and Kartik Sheth for commenting on the manuscript. This work is
based on observations and archival data made with the Spitzer Space Telescope,
which is operated by the Jet Propulsion Laboratory, California Institute of
Technology under a contract with NASA. We are grateful to the dedicated staff
at the Spitzer Science Center for their help and support in planning and execu-
tion of this Exploration Science program. We also gratefully acknowledge sup-
port from NASA JPL/Spitzer grant RSA 1374189 provided for the S4G project.
We acknowledge financial support to the DAGAL network from the People
Programme (Marie Curie Actions) of the European Union’s Seventh Framework
Programme FP7/2007-2013/ under REA grant agreement number PITN-GA-
2011-289313. This research has made use of the NASA/IPAC Extragalactic
Database (NED) which is operated by the Jet Propulsion Laboratory, California
Institute of Technology, under contract with the National Aeronautics and Space
Administration.
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Appendix A: Fits of the radial luminosity profiles
This Appendix presents the radial fits to the luminosity profiles of the galaxies
with a detectable CMC. The fits are displayed as in Fig. 1.
10
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
0 20 40 60
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
ESO 157-49
-60 -40 -20 0 20 40 60
-20
-10
0
10
20
0 50 100 150 200 250
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
ESO 240-11
-200 -100 0 100 200
-20
0
20
0 10 20 30 40 50 60
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
ESO 346-1
-60 -40 -20 0 20 40 60
-10
-5
0
5
10
0 50 100 150
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
ESO 440-27
-100 0 100
-40
-20
0
20
0 20 40 60 80
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
ESO 443-21
-50 0 50
-20
-10
0
10
20
0 20 40 60 80
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
ESO 533-4
-50 0 50
-20
-10
0
10
0 10 20 30 40 50 60
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
ESO 544-27
-60 -40 -20 0 20 40 60
-10
0
10
0 20 40 60 80
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
IC 610
-50 0 50
-20
-10
0
10
20
11
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
0 10 20 30 40 50
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
IC 1553
-40 -20 0 20 40
-20
-10
0
10
20
0 20 40 60 80 100
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
IC 1711
-100 -50 0 50 100
-30
-20
-10
0
10
20
0 50 100 150
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
IC 5176
-100 0 100
-40
-20
0
20
0 10 20 30 40 50
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 489
-40 -20 0 20 40
-20
-10
0
10
20
0 20 40 60 80 100
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 522
-100 -50 0 50 100
-20
-10
0
10
20
0 50 100 150
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 678
-150 -100 -50 0 50 100 150
-40
-20
0
20
40
0 50 100 150
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 1032
-100 0 100
-60
-40
-20
0
20
40
60
0 20 40 60 80 100
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 1163
-100 -50 0 50 100
-20
-10
0
10
20
12
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
0 50 100 150 200
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 1596
-200 -100 0 100 200
-50
0
50
0 20 40 60 80 100 120
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 2732
-100 -50 0 50 100
-20
0
20
0 20 40 60 80
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 3098
-50 0 50
-30
-20
-10
0
10
20
30
0 20 40 60 80 100 120
R (arcsec)
28
27
26
25
24
23
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 3454
-100 -50 0 50 100
-20
-10
0
10
20
0 20 40 60 80 100 120 140
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 3501
-150 -100 -50 0 50 100
-20
-10
0
10
20
0 20 40 60 80
R (arcsec)
28
27
26
25
24
23
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 3592
-50 0 50
-20
-10
0
10
20
0 20 40 60 80 100 120 140
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 3600
-100 -50 0 50 100
-40
-20
0
20
0 100 200 300 400 500
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 3628
-400 -200 0 200 400
-100
-50
0
50
100
13
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
0 50 100 150
R (arcsec)
28
26
24
22
20
18
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 4111
-150 -100 -50 0 50 100 150
-40
-20
0
20
40
0 100 200 300 400
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 4437
-400 -200 0 200
-50
0
50
0 100 200 300 400 500
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 4565
-400 -200 0 200 400
-100
-50
0
50
0 20 40 60 80 100 120 140
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 4607
-100 -50 0 50 100
-30
-20
-10
0
10
20
0 50 100 150 200
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 5529
-200 -100 0 100 200
-20
0
20
0 20 40 60 80 100 120
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 5981
-100 -50 0 50 100
-20
-10
0
10
20
0 20 40 60 80 100 120
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 6010
-100 -50 0 50 100
-40
-20
0
20
0 10 20 30 40 50 60
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
NGC 7347
-40 -20 0 20 40
-20
-10
0
10
14
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
0 20 40 60 80 100 120
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
PGC 13646
-100 -50 0 50 100
-20
-10
0
10
20
0 20 40 60
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
PGC 28308
-60 -40 -20 0 20 40 60
-20
-10
0
10
20
0 10 20 30 40 50 60
R (arcsec)
28
27
26
25
24
23
22
µ 
(m
ag
 a
rc
se
c-
2
)
PGC 32548
-60 -40 -20 0 20 40
-15
-10
-5
0
5
10
0 20 40 60 80 100 120
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
PGC 52809
-100 -50 0 50 100
-30
-20
-10
0
10
20
0 20 40 60 80
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 903
-50 0 50
-30
-20
-10
0
10
20
30
0 10 20 30 40 50
R (arcsec)
28
27
26
25
24
23
22
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 5347
-40 -20 0 20 40
-10
-5
0
5
10
0 10 20 30 40 50 60
R (arcsec)
28
27
26
25
24
23
22
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 5689
-60 -40 -20 0 20 40 60
-20
-10
0
10
0 20 40 60 80
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 7086
-50 0 50
-30
-20
-10
0
10
20
15
Comero´n, S., et al.: Evidence for the concurrent growth of thick discs and central mass concentrations from S4G imaging
0 20 40 60 80 100 120
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 8737
-100 -50 0 50 100
-20
-10
0
10
20
0 20 40 60 80
R (arcsec)
28
27
26
25
24
23
22
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 9448
-50 0 50
-10
0
10
0 20 40 60 80
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 9665
-50 0 50
-20
-10
0
10
20
0 20 40 60 80 100
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 10043
-100 -50 0 50
-10
0
10
0 50 100 150 200
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 10288
-200 -100 0 100 200
-20
0
20
0 20 40 60
R (arcsec)
28
26
24
22
20
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 12518
-60 -40 -20 0 20 40 60
-20
-10
0
10
0 20 40 60 80
R (arcsec)
28
26
24
22
µ 
(m
ag
 a
rc
se
c-
2
)
UGC 12857
-50 0 50
-20
-10
0
10
20
16