11institutetext: European Space Astronomy Centre, ESA, PO Box 78, 28691
Villanueva de la Cañada, Madrid, Spain 22institutetext: Telespazio Vega UK
Ltd. for ESA/ESAC, Madrid, Spain 33institutetext: Observatorio Astronómico
Nacional (IGN), Ap 112, 28803 Alcalá de Henares, Spain 44institutetext:
Observatorio Astronómico Nacional (OAN-IGN), Calle Alfonso XII 3, 28014
Madrid, Spain

# Fast outflows in protoplanetary nebulae and young planetary nebulae observed
by Herschel/HIFI

M. Lorenzo currently at Centro de Astrobiología (CSIC-INTA), Torrejón de
Ardoz, Madrid, Spain11 D. Teyssier 22 V. Bujarrabal 33 P. García-Lario 11 J.
Alcolea 44 E. Verdugo 11 A. Marston 11

###### Abstract

Context. Fast outflows and their interaction with slow shells (generally known
as the fossil circumstellar envelope of asymptotic giant branch stars) play an
important role in the structure and kinematics of protoplanetary and planetary
nebulae (pPNe, PNe). To properly study their effects within these objects, we
also need to observe the intermediate-temperature gas, which is only
detectable in the far-infrared (FIR) and submillimetre (submm) transitions.

Aims. We study the physical conditions of the outflows presented in a number
of pPNe and PNe, with a focus on their temperature and excitation states.

Methods. We carried out Herschel/HIFI observations in the submm lines of 12CO
in nine pPNe and nine PNe and complemented them with low-$J$ CO spectra
obtained with the IRAM 30m telescope and taken from the literature. The
spectral resolution of HIFI allows us to identify and measure the different
nebular components in the line profiles. The comparison with large velocity
gradient (LVG) model predictions was used to estimate the physical conditions
of the warm gas in the nebulae, such as excitation conditions, temperature,
and density.

Results. We found high kinetic temperatures for the fast winds of pPNe,
typically reaching between 75 K and 200 K. In contrast, the high-velocity gas
in the sampled PNe is colder, with characteristic temperatures between 25 K
and 75 K, and it is found in a lower excitation state. We interpret this
correlation of the kinetic temperature and excitation state of fast outflows
with the amount of time elapsed since their acceleration (probably driven by
shocks) as a consequence of the cooling that occurred during the pPN phase.

###### Key Words.:

stars: AGB & post-AGB - circumstellar matter - stars:mass-loss - planetary
nebulae: general

##  1 Introduction

Throughout the very short post-AGB phase, lasting less than a few thousand
years, a dramatic change occurs in the morphology and dynamics of the
circumstellar envelopes (CSEs) of evolved stars of intermediate mass. Starting
from an intense mass-loss period that results in the formation of a dense,
spherical, and slowly expanding CSE, asymptotic giant branch (AGB) stars then
turn into the central stars of planetary nebulae (PNe). These PNe present
complex morphological structures, often with strong axial-symmetry. These
extreme kinetic and morphological changes are thought to arise from the
interaction between fast, collimated (jet-like) winds and the slow, dense
layers of the CSE. The transformation takes place in the intermediate stage
known as the protoplanetary nebula (pPN) phase.

Collimated winds, commonly known as fast outflows, are ejected close to the
central core and carry a significant fraction of nebular mass, that is, $\sim$
0.1-0.3 M⊙ (Bujarrabal et al. 2001), with velocities around 100 km s-1. The
interaction between these fast outflows and the remnant of the slowly
expanding circumstellar envelope (AGB-CSE) is believed to be the main agent
responsible for the change of kinematics and structure observed in pPNe,
sculpting the final shape of PNe (Balick & Frank 2002). Several observations
of molecular lines have identified fast outflows in many pPNe and PNe; see
examples in Bujarrabal et al. 1998, 2001; Alcolea et al. 2001; Castro-Carrizo
et al. 2002; Sánchez Contreras et al. 2004; Fong et al. 2006, and references
therein. In most cases, they exhibit strong emission of molecular lines with
composite profiles. Typically, these profiles consist of a central core that
comes from the earlier slow AGB wind and broad line wings arising from fast
post-AGB outflows. In molecular line emission, these outflows show bipolar or
hourglass-like geometries and dynamics in almost all well-studied cases,
although the optical images of many PNe and pPNe tend to be more complex, with
multipolar or point-like symmetries.

Until now, most of what we know about the structure, dynamics, and physical
conditions of molecular material in nebulae has been derived mainly from
observations of low-$J$ transitions of molecular lines, especially the $J$=2–1
and $J$=1–0 transitions of CO. These lines primarily trace molecular gas in
the cold, (typically) outer layers of the circumstellar envelope, where the
expansion velocity is already constant and the temperature, as it is
decreasing with the radius as $\sim$ 1/r, typically reaches between 10-20 K.
There are also many studies analysing molecules with higher dipole moments,
such as HCN, HCO+, and HNC, which trace denser gas than CO (e.g. Bachiller et
al. 1997; Zhang et al. 2008; Milam et al. 2009; Edwards & Ziurys 2013; Edwards
et al. 2014; Schmidt & Ziurys 2016; Bublitz et al. 2019, and references
therein). However, to completely understand the shaping mechanism of the
highly collimated jets in pPNe and PNe, and, consequently, their evolution, we
also need to observe warmer regions (100 K $\lesssim$ TK $\lesssim$ 1000 K),
which is where acceleration takes place. A proper study of intermediate-
temperature gas requires observations of higher $J$ transitions of CO that are
found in the far-infrared (FIR) and submillimetre (submm) ranges.

Unfortunately, there are very few studies based on the observation of the
molecular gas of pPNe and PNe in the FIR and sub-mm lines. Justtanont et al.
(2000) were the first to analyse the intermediate-temperature gas around
evolved stars in the FIR. They presented ISO-LWS observations of high-$J$
transitions of CO (from $J$=14–13 up to $J$=37–36) in CRL 618, CRL 2688 and
NGC 7027, but the low spectral resolution of their data prevented them from
identifying the nebular components responsible for the line emission. Higher
spectral resolution was achieved with Herschel in the FIR/submm observations
presented in Bujarrabal et al. (2012). By comparing their results with
theoretical predictions, they found warm fast winds in CRL 618 and CRL 2688
(Tk $\geq$ 100 K) and colder outflows in OH 231+4.2 and NGC 6302 (Tk $\sim$ 30
K).

Herschel, and in particular its HIFI heterodyne spectrometer, provides the
necessarily high spectral resolution ($\frac{\Delta\upsilon}{\upsilon}$ ¡
10-6) in the FIR and submm to determine the dynamics and physical conditions
of the shocked layers of this type of nebulae and to resolve the different
features of their profiles. In this paper, we present observations of
intermediate-excitation CO lines in a sample of eighteen nebulae carried out
with HIFI. These observations have been obtained as part of the guaranteed-
time project SUbmilimetre Catalogue of Circumstellar Envelopes of StarS
(SUCCESS), a submm/FIR CO survey of a large sample of AGB and post-AGB stars.
In particular, this programme obtained observations of the 12CO $J$=5–4 and
$J$=9–8 lines in 18 post-AGB stars. Of this sample of nebulae, three were also
observed in the 12CO $J$=14–13 line and five were followed up in 12CO $J$=2–1
and $J$=1–0 lines using the IRAM 30m telescope.

##  2 Observations

###  2.1 Herschel/HIFI observations

Observations in the submm lines of CO were obtained in the context of the
Herschel SUCCESS Guaranteed Time project (Teyssier et al. 2011) executed with
the HIFI high-resolution spectrometer (de Graauw et al. 2010). The post-AGB
stars considered in this paper represent about a quarter of the whole sample
observed by this programme. Details on how the data were acquired and
processed can be found in Danilovich et al. (2015). No particular spectral
smoothing was applied to the data for their analysis, however, rebinning was
applied in some of the spectra given in the paper.

###  2.2 Supplementary data

Our HIFI data were complemented with low-$J$ CO spectra obtained during an
observing campaign at the IRAM 30m telescope, aimed at getting better
calibrated data for some of the sources where the only available spectra had
been taken several decades ago. These observations were conducted in December
2012 and comprised five sources from our sample of post-AGB stars. For the
rest of the sources, we relied on data digitised from the literature or
unpublished data available to members of our team. Table 1 summarises the
origin of the spectra used in this paper for the low-excitation CO lines. The
corresponding spectra are shown in Appendix A. In the analysis described in
Section 4, we distinguish among the compact and extended sources and we
provide examples of the information shown together with the spectra of these
sources in Figures 1 and 2, respectively. For some sources, spectra are also
given for additional CO transitions (e.g. Fig.14). These were used in support
to the selection of the velocity ranges applied in our analysis (Section 4.1).

In two cases (IRAS 21282+5050 and OH17.7-2.0), the CO data stem from
interferometric observations as no single-dish counterpart could be found in
the literature. The two sources concerned by these observations are either
compact (OH17.7-2.0), or their interferometric data were combined with short
spacings obtained at the 30m (IRAS 21282+5050). We have therefore assumed that
there is no significant flux loss in the corresponding data. The conversion
from surface brightness to temperature units we used were: 45 mK Jy-1 beam-1
(CO $J$=1–0) and 29 mK Jy-1 beam-1 (CO $J$=2–1) for the IRAM Plateau de Bure
observations of Castro-Carrizo et al. (2010), and 1.19 K Jy-1 beam-1 (CO
$J$=1–0) for the OVRO observations of Sánchez Contreras et al. (2007).

Table 1: Details on the origin of the low-$J$ CO spectra used in this study.

Source | CO $J$=1–0 | CO $J$=2–1
---|---|---
IRAS 07134+1005 |  Bujarrabal et al. (1992)a
IRAS 19500–1709 | Bujarrabal private communicationa
IRAS 21282+5050 |  Castro-Carrizo et al. (2010)a,b
IRAS 22272+5435 | Bujarrabal private communicationa
M 1–92 | New 30m data
M 2–56 | Bujarrabal private communicationa
OH 17.7–2.0 |  Sánchez Contreras et al. (2007)c |  Heske et al. (1990)a
R Sct | Bujarrabal private communicationa
89 Her | Bujarrabal private communicationa
NGC 6537 | New 30m data
M 1–16 | New 30m data
M 1–17 | New 30m data
M 3–28 | New 30m data
NGC 2346 |  Bachiller et al. (1989)a
NGC 6072 |  Cox et al. (1992)d
NGC 6720 |  Bachiller et al. (1989)a
NGC 6781 |  Bachiller et al. (1993)a
NGC 7293 |  Zack & Ziurys (2013a)e |  Young et al. (1999)f

  * •

Notes. a IRAM 30 m telescope, b IRAM Plateau de Bure interferometer, c OVRO
interferometer, d SEST telescope, e ARO 12 m telescope, f CSO 10.4 m
telescope.

Table 2: Frequencies and telescope parameters for the CO transitions used in
this work.

Transition |  | Frequency
---
(GHz)
Telescope |  | $\theta$
---
(”)
12CO $J$=1–0 | 115.271 | IRAM | 21.4
|  | PdBI | 4.9$\times$3.9
|  | OVRO | 10.9$\times$7.1
|  | SEST | 45
|  | ARO | 55
12CO $J$=2–1 | 230.538 | IRAM | 10.7
|  | PdBI | 2.6$\times$1.4
|  | SEST | 23
|  | CSO | 31
12CO $J$=3–2 | 345.796 | APEX | 17.3
|  | ARO SMT | 22
12CO $J$=4–3 | 461.041 | ARO SMT | 16
12CO $J$=5–4 | 576.268 | HIFI | 36.1
12CO $J$=6–5 | 691.473 | HIFI | 30.4
|  | ARO SMT | 11
12CO $J$=9–8 | 1036.912 | HIFI | 20.1
12CO $J$=10–9 | 1151.985 | HIFI | 18.2
12CO $J$=14–13 | 1611.794 | HIFI | 12.9
12CO $J$=16–15 | 11841.345 | HIFI | 11.5

  * •

Notes. IRAM is the 30 m telescope at the Institute de Radioastronomie
Millimétrique; PdBI refers to the IRAM Plateau de Bure interferometer (now
known as NOEMA); OVRO refers to the millimeter interferometer of the Owens
Valley Radio Observatory; SEST is the 15 m Swedish-ESO Sub-millimeter
Telescope; ARO refers to the 12 m telescope of the Arizona Radio Observatory;
CSO is the 10.4 m Leighton telescope of the Caltech Submilimeter Observatory;
ARO SMT is the 10 m Submillimeter Telescope of the Arizona Radio Observatory;
APEX refers to the Atacama Pathfinder EXperiment; and HIFI is the Heterodyne
Instrument for the Far-Infrared aboard Herschel.

##  3 Description of the sample

Table 3: Observed sources.

Source |  | $\alpha$
---
(J2000)
| $\delta$
---
(J2000)
Type | SpT |  | $d$
---
(kpc)
| $L$
---
(103 L⊙)
| $v_{\rm sys}$
---
(km s-1)
| $\dot{M}$
---
(M⊙ yr-1)
Morph. | Ref.
pPNe |  |  |  |  |  |  |  |  |  |
IRAS 07134+1005 | 07 16 10.258 | +09 59 47.99 | C | F5Iab | 3 | 13.5 | 72 | $5\times 10^{-5}$ | b | 1, 2
IRAS 19500–1709 | 19 52 52.701 | -17 01 50.30 | C | F2-6 | 1 | 1.5 | 25 | - | - | 1
IRAS 21282+5050 | 21 29 58.420 | +51 03 59.80 | C | 09.5, WC11 | 3 | 5.3 | 14 | 6$\times$10-5 | - | 1, 2
IRAS 22272+5435 | 22 29 10.374 | +54 51 06.35 | C | G5Ia | 1.7 | 8.3 | -28 | 2$\times$10-5 | s | 1
M 1–92 | 19 36 18.910 | +29 32 50.00 | O | B0.5IV | 2.5 | 10 | -1 | 2.2$\times$10-4 | b | 1, 3
M 2–56 | 23 56 36.400 | +70 48 18.20 | - | Be | 3 | 10 | -27 | 2.3$\times$10-5 | b | 1
OH 17.7–2.0 | 08 30 31.000 | -14 28 57.00 | O | F0 | 2 | 2.9 | 61 | 1$\times$10-4 | b | 1, 4
R Sct† | 18 47 28.950 | -05 42 18.53 | O | G0-K2 | 0.4 | 4 | 56 | 3.3$\times$10-7 | - | 1
89 Her† | 17 55 25.189 | +26 02 59.97 | O | F2Ibe | 0.6 | 3.3 | -8 | 4$\times$10-6 | b | 1, 5
Compact PNe |  |  |  |  |  |  |  |  |  |
NGC 6537 | 18 05 13.104 | -19 50 34.88 | O* | - | 1.6 | 2 | 8.6 | 1$\times$10-7 | b | 6, 7, 8
M 1–16 | 07 37 18.955 | -09 38 49.67 | O | - | 1.8 | 0.12 | 50 | 2$\times$10-5 | e | 2
M 1–17 | 07 40 22.206 | -11 32 29.81 | - | - | 5.8 | - | 28 | - | c | 9, 10
M 3–28 | 18 32 41.288 | -10 05 50.03 | C | - | 2.5 | - | 32 | - | q | 6, 11
Extended PNe |  |  |  |  |  |  |  |  |  |
NGC 2346 | 07 09 22.521 | -00 48 23.60 | - | A5V | 0.9 | - | 2 | - | b | 6
NGC 6072 | 16 12 58.079 | -36 13 46.06 | - | B | 1.39 | 0.39 | 15 | - | b | 6
NGC 6720 | 18 53 35.079 | +33 01 45.03 | O | hgO(H) | 0.7 | 1 | -4 | 2$\times$10-4 | b | 6, 2
NGC 6781 | 19 18 28.085 | +06 32 19.29 | C | DAO | 0.95 | 0.35 | 17 | - | e | 6
NGC 7293 | 22 29 38.550 | -20 50 13.60 | C | DAO.5 | 0.22 | 0.09 | -24 | 1$\times$10-4 | b | 6, 2

  * •

†: Bipolar nebula with rotational disk

  * •

*: Edwards & Ziurys (2013) questioned how O-rich NGC 6537 is.

  * •

Morph. abbreviations: b: bipolar or hourglass-like, c:compact, e: elliptical,
q: quadrupolar, s: spherical.

  * •

References: (1): Bujarrabal et al. (2001); (2): Fong et al. (2006); (3):
Bujarrabal et al. (1998); (4): Bujarrabal et al. (1994); (5): Alcolea &
Bujarrabal (1995); (6): Frew (2008); (7): Cuesta et al. (1995); (8): Edwards &
Ziurys (2013); (9): Maciel (1984); (10): Bachiller et al. (1991); (11): Ziurys
et al. (2020).

Figure 1: Observations of IRAS 21282+5050. Shadowed regions mark the LSR
velocity ranges adopted to determine the characteristic intensity of each
feature. The dashed vertical lines indicates the systemic velocity as
tabulated in Table 3. Figure 2: Observations of NGC 6720. The shadowed regions
indicate the velocity ranges adopted to determine the characteristic intensity
of the detected velocity components; see Section 4.1 for further details.

Our sample is presented in Table 3, alongside with their observed coordinates
and other parameters. It consists of 18 post-AGBs: nine protoplanetary nebulae
and nine planetary nebulae. Five of the PNe are considered extended sources
with respect to the size of the HIFI beams (see their values in Table 2). In
this section, we present a brief description of each source and a qualitative
analysis of their observed CO profiles.

###  3.1 IRAS 07134+1005

IRAS 07134+1005 (HD 56126) is a pPN with a double-peaked SED (Kwok et al.
1989), presenting a high abundance of s-process elements, a typical
characteristic of post-AGBs that have experienced the third dredge-up. Mid-IR
images of IRAS 07134+1005 (Meixner et al. 1997) show an elliptical detached
shell surrounding two aligned peaks, which these authors interpreted as limb-
brightened peaks of a shell with an equatorial density enhancement. The
profile of the CO $J$=9–8 line obtained in our observations shows a double
peak around the line core, indicative of high-velocity shocks (Fig. 6).

###  3.2 IRAS 19500–1709

IRAS 19500–1709 (HD 187885) is a carbon-rich pPN. The high abundance of
s-process elements in its shell and a feature at 21 $\mu$m in its spectra (Van
Winckel & Reyniers 2000) suggests that the third dredge-up has already
occurred in this object. IRAS 19500–1709 has a bi-modal spectral energy
distribution (SED) characteristic of pPN with detached circumstellar dust
envelope (Hrivnak et al. 1989). Our observed lines reflect this composite
structure. They show, in most of them, a central line core (with an expansion
velocity of 11 km s-1) and wide line wings whose expansion velocity goes up to
$\sim$ 30 km s-1. Figure 7 shows that the contribution from the remnant of the
slow AGB wind is dominating the profile at low $J$, but then decreases
significantly in the $J$=5–4 line and above. This suggests that the
temperature of the fast outflows is high in contrast with the slow expanding
AGB-CSE. In addition, we can see a clear absorption in the blue wing of
$J$=2–1 line, supporting the idea that the fast component presents higher
temperatures than the slow gas and that it is found closer to the central
star.

###  3.3 IRAS 21282+5050

IRAS 21282+5050 is a young, C-rich protoplanetary nebula. The CO $J$=1–0 maps
by Shibata et al. (1989) show a clumpy expanding shell whose nature was
interpreted as an expanding toroidal structure with an axis along the east-
west direction, perpendicular to the line of sight. The spectra of the lower-
energy transitions displayed in Figure 1 exhibit profiles dominated by a
roughly symmetric component around the systemic velocity. However, as the
energy of the transition increases, the profile seems more of a composite as
well as steep-sided, revealing the three typical features: two line wings
growing in intensity with increasing $J$ and a line-core decreasing with the
excitation, indicative of a high temperature in the object’s outflows.

###  3.4 IRAS 22272+5435

IRAS 22272+5435 (HD 235858) is a C-rich pPN (Kwok et al. 1989; Omont et al.
1995). Studies at optical (Ueta et al. 2000) and near-infrared (Gledhill et
al. 2001) wavelengths have identified an axisymmetric nature (P.A. $\sim$
-35º) for the nebula and a toroidal dust shell normal to the axis. Maps of CO
$J$=1–0 (Fong et al. 2006) have shown that its molecular gas expands following
a spherical structure with multiple protrusions in a quadrupolar geometry. As
can be seen in Fig. 8, its profile displays a parabolic-like shape as commonly
observed in AGB profiles.

###  3.5 M 1–92

M 1–92, also known as Minkowski’s Footprint, is a very well studied O-rich pPN
that presents OH masers. It shows a clear hourglass-like structure consisting
of a thin flat disk expanding radially and two symmetrical lobes (Alcolea et
al. 2007). The CO maps show that most of its molecular gas is located in the
equatorial disk and in the walls of two empty lobules, where high velocities
are observed (Bujarrabal et al. 1998). Figure 9 shows intermediate velocity
features that disappear as the energy levels increase.

###  3.6 M 2–56

M 2–56 is a pPN with a bipolar appearance in the optical and an hourglass-like
shape in CO maps. These maps show a dense ring probably arising from the
fossil AGB envelope and two high-velocity components at 11”-15” from the
centre along the symmetry axis. It also presents two empty attached shells
close to the equatorial disk, expanding along the nebula axis at low
velocities (Castro-Carrizo et al. 2002). As we can seen in Fig. 10, the CO
lines of M 2–56 are composite with a strong blue wing that remains as the
dominant feature of the spectral profile at high values of $J$. This
remarkable difference between the blue and the red wings is probably due to an
intrinsic asymmetry.

###  3.7 OH 17.7–2.0

OH17.7-2.0 (IRAS 18276+143) is an OH/IR pPN that presents a double-peaked SED
indicative of a detached dust shell (Le Bertre 1987). Polarised intensity
images (Gledhill 2005) have found hourglass-like geometry with a symmetric
axis at P.A. $\sim$ 23º. Horn-like features seen in its position-velocity
diagram suggest the presence of outflows close to the source (Bains et al.
2003). A relative absence of OH masers is observed in its polar regions,
compared with the strong emission detected near its equator (also seen in M
1–92). This fact suggests that these areas may have been excavated by fast
outflows creating empty symmetric lobes of gas (Sánchez Contreras et al.
2007). Although our HIFI observations for this source are relatively noisy
(Fig. 11), for CO $J$=9–8, they reveal in a feature around 40 km s-1 expanding
at $\sim$ 10 km s-1.

###  3.8 R Sct

R Sct (HR 7066, HIP 92202) is a peculiar RV Tau variable whose evolutionary
stage is still under debate. It presents a complex profile in all CO lines
(Fig 12). In the low-$J$ transition $J$=2–1, we observe a narrow double peak,
which is a characteristic property of rotating disks (Bujarrabal et al. 2013).
A strong blue wing is seen at these low energy levels, but as $J$ increases,
the emission of this feature is abated, suggesting a low temperature for the
approaching outflow. Within our sample, this source is the one exhibiting the
narrowest line profile ($\sim$ 10 km s-1).

###  3.9 89 Her

89 Her (HR 6685) is a low-mass pPN. It is a close binary system known to
present a remarkable near-infrared excess due to the emission of hot dust at
$\sim$ 1000 K (Waters et al. 1993; de Ruyter et al. 2006) and a clear far-
infrared excess suggesting the presence of large grains in its dusty envelope
(Shenton et al. 1995). In the $J$=2–1 and $J$=1–0 CO lines presented in Fig.
13, we observe sharp profiles with a prominent central peak and weak wings at
low velocities. At higher $J$, the structure of the profile does not change,
yet their emission becomes weaker and approaches the noise level, while the
wings are practically undetectable. According to its CO maps (Bujarrabal et
al. 2007, 2013), this profile corresponds to a compact and unresolved
component in the centre of the nebula which is suspected to be a rotating disk
and an hourglass-like extended outflow (with a maximum velocity of $\sim$ 7 km
s-1).

###  3.10 NGC 6537

The Red Spider nebula is one of the PNe with the highest excitation known thus
far (Ashley & Hyland 1988). Its spider morphology is mainly observed in the
atomic emission lines, but it can also be seen in molecular hydrogen. NGC 6537
is hourglass-like and possibly multi-polar in its nebular shape (Cuesta et al.
1995). Because this PN is relatively young (its age is estimated to be $\sim$
1600 yr, Matsuura et al. (2005)), it has been suggested that the remnant
molecular gas has not yet been accelerated enough in the polar outflows
(Edwards et al. 2014). In Fig. 14, we observe that the strong blue wing, seen
as a left peak in $J$=2–1 and $J$=1–0 lines, decreases with $J$, suggesting a
low temperature in the approaching outflow. There is an important interstellar
medium contamination between 20 and $\sim$ 45 km s-1 in the two lowest-$J$ CO
lines, which we have tried to avoid in our study. The same tendency is seen in
$J$=6–5 and $J$=4–3 lines presented in Edwards et al. (2014).

Although the FIR emission of this source fills the HIFI beam for the $J$=5–4
line entirely and exceeds it for the other lines, we considered it as a
compact PN because we pointed at its centre and not their external velocity
components, in contrast with what was carried out for the extended PNe.
However, the fact that the lines do not come from an area smaller than the
HIFI beams prevent us from performing accurate comparisons with the other
compact objects and theoretical models (see Section 4).

###  3.11 M 1–16

M 1–16 is a young, compact planetary nebula known to exhibit multiple outflows
with different velocities. It presents a slow outflow (with an expansion
velocity of 19 km s-1), which exhibits an elliptical structure and high-
velocity flows consisting of three bipolar jets with shock velocities of
$\sim$ 300 km s-1. These fast outflows have excavated extended bipolar
cavities producing symmetric lobes of molecular gas (Huggins et al. 2000). The
most remarkable characteristics seen in our observed lines presented in Fig.
15 are the extended and strong wings, indicative of high-velocity outflows and
their decreasing contribution with the increase of the excitation, leaving a
peak with a broad but weak pedestal in $J$=9–8.

###  3.12 M 1–17

M 1–17 is a compact planetary nebula whose emission is weak at practically all
wavelengths, except at the observed CO lines. Bachiller et al. (1991) studied
this source in $J$=2–1 and $J$=1–0 and identified two components in its
profile: a slow wind, with a double-peaked profile extending over 34 km s-1
and a fast wind, with a triangular-shape profile extending over 78 km s-1. Our
observations at intermediate-energy transitions (Fig. 16) provide additional
information about the molecular content of this nebula and reveal in
particular a new component further away from the systemic velocity, expanding
over less than 5 km s-1. As the energy increases, the medium-velocity
components known so far lose intensity and the component in the outskirts
kicks in, indicating a higher temperature in the latter and hinting towards a
possible acceleration gradient.

###  3.13 M 3–28

M 3–28 is a small PN, with a bipolar, butterfly-like morphology and a
filamentary structure (Calvet & Cohen 1978). As shown in Fig. 17, the line
profile is relatively broad, but its exact extent at the red end cannot be
accurately assessed because of the strong interstellar medium contamination
that is especially visible in the low excitation lines. Our observation at
higher $J$, however, allows us to appreciate, for the first time, the full
breadth of the profile, which spans nearly 80 km s-1 overall. This line
contamination was taken into account when selecting the velocity ranges of
interest for this study (see Section 4.1). Apart from that, the global profile
is that of a double peak on either side of a weak line core, similar to that
observed in other sources, such as M 1–17 or IRAS 21282+5050.

###  3.14 NGC 2346

NGC 2346 is an extended bipolar PN. It has an hourglass-like geometry
consisting of two open bipolar lobes and an expanding toroidal structure in
its centre. This equatorial torus consists of large amounts of molecular gas
and dust, and it is tilted 56∘ with respect to the plane of the sky (Su et al.
2004). The source is extended compared to the beams of our observations (see
Fig. 47) and the targeted region corresponds to an edge of this toroidal
structure. Because, in terms of size, the CO $J$=2–1 beam is smaller than the
others, our profile consists of just one feature at the systemic velocity of
the nebula. In contrast, in the other CO lines, we observe four features
corresponding to four different velocity components (Fig 18). Their emission
drastically decreases at high-energy transitions, suggesting a low temperature
in all these components.

###  3.15 NGC 6072

Near-infrared and molecular hydrogen images (Kwok et al. 2010) have
characterised NGC 6072 as having an elliptical, equatorial ring and three
pairs of lobes, indicating a multipolar structure. This morphology is the
result of fast outflows interacting with the slow winds of the AGB phase and
whose orientation is changing over time. Most of our observed CO lines present
a double-peaked profile, indicative of a detached molecular shell expanding at
$\sim$ 15 km s-1 (Fig. 19). However, the emission is fully extinguished in the
CO $J$=14–13 line, suggesting a low temperature in this expanding shell.

###  3.16 NGC 6720

The Ring Nebula is an extended, O-rich planetary nebula with a bipolar
morphology tilted with respect to the line of sight (Kwok et al. 2008; O’Dell
et al. 2013). Observations of CO $J$=2–1 and $J$=1–0 lines carried out by
Bachiller et al. (1989) and presented in Fig. 2 revealed unusually-high-
velocity molecular gas. Owing, most likely, to the larger size of the CO
$J$=5–4 beam in our HIFI observations, we detected two more velocity
components at roughly $\sim$ 15 km s-1 on either side of the systemic
velocity. The fact that the emission of all of these components remains
visible in the CO $J$=9–8 line suggests high-temperature molecular gas at the
outskirts of this nebula.

###  3.17 NGC 6781

NGC 6781 has an elliptical morphology very similar to that of the Ring nebula.
High-resolution mapping of CO $J$=2–1 and $J$=1–0 lines by Bachiller et al.
(1993) has provided a detailed kinematic structure of its envelope. A model,
consisting of a thin, expanding, ellipsoidal shell, truncated at the two ends,
was proposed by the authors to explain the observation. As we can see in Fig.
20, our line profiles are very complex, revealing a clumpy nature for the CO
shell.

###  3.18 NGC 7293

The Helix nebula is the closest planetary nebula of our sample. It presents an
extreme bipolar morphology tilted with respect to the line of sight and a
central toroidal ring. This last feature expands around the ionized nebula
with arcs and filaments at high latitudes (Meaburn et al. 2005; Zeigler et al.
2013). The edge of one of these filaments corresponds to the position we have
observed in this work, at offset (-435”, 75”) from the central star (Fig. 47).
As can be seen in Fig. 21, the profile reveals various features that we treat
as separate velocity components resolved along the line of sight. Their
relative intensity remains constant, except in the CO $J$=9–8 line, where
their emission drops drastically, indicating that none of the observed
components exhibits a high temperature. This observation suggests that most of
the molecular gas of the PN has not recently been ejected, an assumption that
is also supported by the findings of Zack & Ziurys (2013b) and Schmidt et al.
(2018).

##  4 Analysis and discussion

###  4.1 Characteristic intensities

The different velocity components evidenced in our high-resolution profiles of
the CO lines provide detailed information about the excitation condition in
the outflows of our objects. In particular, the relative strength of the
respective spectral components and their evolution with $J$ can vary
significantly between the various sources and we use this signature here as a
proxy of the evolutionary state of the nebula.

For this, we follow the method presented in Bujarrabal et al. (2012), where
the contribution of each line component is estimated through its
characteristic intensity over a selected velocity range in the relevant CO
transition. We note beforehand that the physical meaning of the selected
velocity ranges is different depending on the relative size of our objects
with respect to the diameter of the instrument beams (for the record, the
applicable beam sizes are collected in Table 2 and plotted on their optical
and FIR images in Appendix B). In the case of compact objects, we interpret
the velocity components as those of a central line core associated with the
remnant of the slow AGB wind, bracketed by two (blue and red) line wings,
assumed to come from the fast outflows resulting from the post-AGB axial
acceleration. For two specific sources (M 1–92 and M 1–17), we also considered
additional intermediate velocity ranges, that we analyse in Section 4.4. In
the case of sources spatially resolved in most of the probed CO line
(thereafter, ‘extended’ sources), we interpret the various features in their
profiles as different velocity components probed along the line of sight
crossing the nebular gas, mostly in their outskirts.

In order to define the respective velocity ranges in a systematic and
objective fashion, we have taken into account the information provided in the
low-$J$ transition of CO, the systemic velocity assumed for the object, and
the spatial distribution of the CO emission available, and for the extended
sources, in particular. We followed a systematic procedure consisting of the
selection of the velocity range of the central core that is consistent with
the systemic velocity of the source. The width of the velocity component was
chosen as a compromise between a sufficient number of spectral channels
compared to the noise level and line ranges being sufficiently separated to
avoid too much mixing between the gas probed in the respective regions. The
same width was then applied to the regions associated to the red and blue
wings. Overall, this leads to red and blue regions being relatively
symmetrical on either side of the central core. For the narrowest profiles,
the requirement to separate the envelope regions probed in each component is
hard to fulfil, but we checked that small adjustments up and down the chosen
widths do not impact the conclusions derived in this paper. In the case of M
3–28, the strong interstellar medium contamination, which is especially
visible in the lower-$J$ transitions was taken into account to adjust the
position of the red wing, avoiding that it would fall too close to this
contribution. For M 1–17 and M 1–92, additional velocity regions were
considered, as explained in Section 4.4.

For the extended source sample, the various features are interpreted as
separate velocity components along the line of sight. For this reason, the
selection essentially considered the three main lines contributing to the
observed profiles. The strategy regarding the choice of the velocity region
width was the same as for the compact sources. In all cases, except for NGC
7293, one of the components is centred exactly on the assumed systemic
velocity. In NGC 7293, the systemic velocity of -24 km s-1 does not match well
with the blue-shifted peak observed in the higher spectral resolution profile
obtained for $J$=2–1 and above. The velocity corresponding to this peak (-21.7
km s-1) was, therefore, assigned instead to one of the components in this
source. In all cases, the choice of the other two components was guided by the
other line peak intensities observed in the line profile, implying that they
did not necessarily fall on either side of the component selected at the
systemic velocity.

The adopted velocity ranges and their characteristic intensities are compiled,
alongside their corresponding 1-$\sigma$ noise uncertainties, in Table 4
(compact objects) and Table 5 (extended objects). In addition, the green,
blue, and red coloured intervals drawn in Figs. 6 to 21 denote the selected
regions.

Table 4: Characteristic intensities of the three representative features
identified in the CO profiles of the compact objects, and their corresponding
coordinates in the pseudo cc-diagram presented in Figs. 3 and 4.

Source |  12CO line |  | Observation
---
ID
| Central core
---
average Tmb(K)
| Blue wing
---
av. Tmb(K)
| Red wing
---
av. Tmb(K)
| Noise
---
r.m.s. (K)
| Line ratios
---
Central core
| Line ratios
---
Averaged line wings
LSR velocity ranges (km s-1) |  | 69.5 – 74.5 | 61.5 – 66.5 | 77.5 – 82.5 |  |  |
IRAS 07134+1005 | $J$=1–0 |  | 1.32 | 0.54 | 1.05 | 0.05 | R54/21 = 0.13 | R54/21 = 0.15
| $J$=2–1 |  | 2.05 | 0.48 | 1.22 | 0.04 |  |
| $J$=5–4 | 1342244408 | 0.263 | 0.090 | 0.134 | 0.004 | R98/54 = 1.07 | R98/54 = 2.05
| $J$=9–8 | 1342229927 | 0.281 | 0.236 | 0.200 | 0.004 |  |
LSR velocity ranges (km s-1) |  | 19.5 – 29.5 | 1 – 11 | 40 – 50 |  |  |
IRAS 19500–1709 | $J$=1–0 |  | 0.83 | 0.15 | 0.06 | 0.01 |  |
| $J$=2–1 |  | 2.06 | 0.62 | 0.35 | 0.01 | R54/21 = 0.05 | R54/21 = 0.11
| $J$=5–4 | 1342216805 | 0.104 | 0.068 | 0.036 | 0.003 |  |
| $J$=9–8 | 1342216811 | 0.224 | 0.162 | 0.076 | 0.003 | R98/54 = 2.16 | R98/54 = 2.24
| $J$=14–13 | 1342231774 | 0.13 | 0.09 | 0.05 | 0.01 |  |
LSR velocity ranges (km s-1) |  | 11.9 – 19.9 | 2.5 – 10.5 | 22 – 30 |  |  |
IRAS 21282+5050 | $J$=1–0 |  | 1.21 | 0.52 | 1.05 | 0.01 |  |
| $J$=2–1 |  | 2.30 | 0.78 | 1.99 | 0.02 | R54/21 = 0.28 | R54/21 = 0.30
| $J$=5–4 | 1342200921 | 0.653 | 0.199 | 0.704 | 0.004 |  |
| $J$=9–8 | 1342210667 | 0.626 | 0.522 | 1.054 | 0.004 | R98/54 = 0.96 | R98/54 = 2.06
| $J$=14–13 | 1342220506 | 0.42 | 0.57 | 0.78 | 0.01 |  |
LSR velocity ranges (km s-1) |  | -30.5 – -25.5 | -37.3 – -32.3 | -23.2 – -18.2 |  |  |
IRAS 22272+5435 | $J$=1–0 |  | 3.14 | 2.03 | 1.69 | 0.02 | R54/21 = 0.034 | R54/21 = 0.03
| $J$=2–1 |  | 8.76 | 3.69 | 3.45 | 0.02 |  |
| $J$=5–4 | 1342200922 | 0.294 | 0.120 | 0.115 | 0.004 | R98/54 = 1.52 | R98/54 = 2.12
| $J$=9–8 | 1342213360 | 0.448 | 0.253 | 0.245 | 0.004 |  |
LSR velocity ranges (km s-1) |  | -7.5 – 7.5 | -41.5 – -26.5 | 26.5 – 41.5 |  |  |
M 1–92 | $J$=1–0 |  | 0.60 | 0.16 | 0.20 | 0.01 | R54/21 = 0.059 | R54/21 = 0.08
| $J$=2–1 |  | 1.48 | 0.29 | 0.55 | 0.02 |  |
| $J$=5–4 | 1342210076 | 0.087 | 0.032 | 0.030 | 0.003 | R98/54 = 0.51 | R98/54 = 1.64
| $J$=9–8 | 1342220500 | 0.045 | 0.045 | 0.055 | 0.002 |  |
LSR velocity ranges (km s-1) |  | -29 – -13 | -52 – -36 | 5 – 21 |  |  |
M 2–56 | $J$=1–0 |  | 0.181 | 0.152 | 0.058 | 0.003 | R54/21 = 0.07 | -
| $J$=2–1 |  | 0.86 | 0.82 | 0.28 | 0.01 |  |
| $J$=5–4 | 1342201534 | 0.058 | 0.067 |  $\lesssim$ 0.005 | 0.003 | R98/54 $\lesssim$ 0.08 | -
| $J$=9–8 | 1342213359 |  $\lesssim$ 0.005 |  $\lesssim$ 0.005 |  $\lesssim$ 0.005 | 0.002 |  |
LSR velocity ranges (km s-1) |  | 59.5 – 64.5 | 49.5 – 54.5 | 69.5 – 74.5 |  |  |
OH 17.7–2.0 | $J$=1–0 |  | 1.48 | 1.13 | 0.84 | 0.07 | R54/21 $\lesssim$ 0.01 | -
| $J$=2–1 |  | 0.98 | 0.48 | 0.44 | 0.05 |  |
| $J$=5–4 | 1342218429 |  $\lesssim$ 0.009 |  $\lesssim$ 0.009 |  $\lesssim$ 0.009 | 0.005 | - | -
| $J$=9–8 | 1342228622 | 0.064 |  $\lesssim$ 0.008 |  $\lesssim$ 0.008 | 0.004 |  |
LSR velocity ranges (km s-1) |  | 54.5 – 56.5 | 51 – 53 | 57.2 – 59.2 |  |  |
R Sct | $J$=1–0 |  | 0.91 | 0.57 | 0.91 | 0.03 | R54/21 = 0.05 | -
| $J$=2–1 |  | 1.48 | 0.53 | 1.36 | 0.02 |  |
| $J$=5–4 | 1342230387 | 0.069 |  $\lesssim$ 0.012 | 0.044 | 0.006 | R98/54 $\lesssim$ 0.16 | -
| $J$=9–8 | 1342229922 |  $\lesssim$ 0.011 |  $\lesssim$ 0.011 |  $\lesssim$ 0.011 | 0.005 |  |
LSR velocity ranges (km s-1) |  | -10 – -6 | -15.5 – -11.5 | -5 – -1 |  |  |
89 Her | $J$=1–0 |  | 0.62 | 0.2 | 0.18 | 0.01 | R54/21 = 0.03 | -
| $J$=2–1 |  | 1.36 | 0.44 | 0.30 | 0.01 |  |
| $J$=5–4 | 1342214335 | 0.043 |  $\lesssim$ 0.010 |  $\lesssim$ 0.010 | 0.005 | R98/54 = 0.71 | -
| $J$=9–8 | 1342218624 | 0.030 |  $\lesssim$ 0.009 |  $\lesssim$ 0.009 | 0.005 |  |
LSR velocity ranges (km s-1) |  | 5.5 – 10.5 | -0.5 – 4.5 | 11.5 – 16.5 |  |  |
NGC 6537 | $J$=1–0 |  | 1.72 | 1.04 | 1.09 | 0.03 |  |
| $J$=2–1 |  | 3.48 | 2.51 | 1.88 | 0.05 | R54/21 = 0.14 | R54/21 = 0.12
| $J$=5–4 | 1342217693 | 0.471 | 0.211 | 0.296 | 0.004 |  |
| $J$=9–8 | 1342215913 | 0.917 | 0.286 | 0.608 | 0.004 | R98/54 = 1.95 | R98/54 = 1.7
| $J$=14–13 | 1342229785 | 0.94 | 0.20 | 0.62 | 0.01 |  |
LSR velocity ranges (km s-1) |  | 45.8 – 54.2 | 35.8 – 44.2 | 55.8 – 64.2 |  |  |
M 1–16 | $J$=1–0 |  | 1.58 | 1.43 | 1.61 | 0.02 | R54/21 = 0.07 | R54/21 = 0.06
| $J$=2–1 |  | 3.91 | 3.24 | 3.74 | 0.03 |  |
| $J$=5–4 | 1342220515 | 0.271 | 0.237 | 0.210 | 0.003 | R98/54 = 0.65 | R98/54 = 0.40
| $J$=9–8 | 1342220495 | 0.176 | 0.083 | 0.097 | 0.003 |  |
LSR velocity ranges (km s-1) |  | 25 – 31 | 3.5 – 9.5 | 47 – 53 |  |  |
M 1–17 | $J$=1–0 |  | 1.16 | 0.09 | 0.13 | 0.03 | R54/21 = 0.06 | R54/21 = 0.13
| $J$=2–1 |  | 1.93 | 0.54 | 0.37 | 0.04 |  |
| $J$=5–4 | 1342218900 | 0.113 | 0.058 | 0.056 | 0.005 | R98/54 = 1.02 | R98/54 = 0.80
| $J$=9–8 | 1342220494 | 0.115 | 0.053 | 0.038 | 0.004 |  |
LSR velocity ranges (km s-1) |  | 29.5 – 34.5 | 15.5 – 20.5 | 41.5 – 46.5 |  |  |
M 3–28 | $J$=1–0 |  | 1.93 | 2.99 | 3.09 | 0.02 | R54/21 = 0.11 | R54/21 = 0.10
| $J$=2–1 |  | 2.61 | 4.22 | 4.90 | 0.04 |  |
| $J$=5–4 | 1342218430 | 0.289 | 0.412 | 0.474 | 0.006 | R98/54 = 0.47 | R98/54 = 0.48
| $J$=9–8 | 1342218214 | 0.136 | 0.174 | 0.257 | 0.004 |  |

Table 5: Characteristic intensities of the identified velocity components and
the coordinates of the CO profiles of extended sources in the pseudo cc-
diagram shown in Fig. 3.

Source |  12CO line |  | Observation
---
ID
| Component #1
---
average Tmb(K)
| Component #2
---
av. Tmb(K)
| Component #3
---
av. Tmb(K)
| Noise
---
r.m.s. (K)
| Averaged line ratio
---
of all components
LSR velocity ranges (km s-1) |  | 12 – 17 | -0.5 – 4.5 | 21.5 – 26.5 |  |
NGC 2346 | $J$=1–0 |  | 0.76 | 0.12 | 0.17 | 0.04 | R54/21 = 0.27
| $J$=2–1 |  | 1.75 | 0.07 | 0.09 | 0.02 |
| $J$=5–4 | 1342218901 | 0.400 | 0.069 | 0.142 | 0.003 | R98/54 $\lesssim$ 0.06
| $J$=9–8 | 1342229928 | 0.061 |  $\lesssim$ 0.007 |  $\lesssim$ 0.007 | 0.004 |
LSR velocity ranges (km s-1) |  | 11 – 19 | -6 – 2 | 25 – 33 |  |
NGC 6072 | $J$=1–0 |  | 0.39 | 0.31 | 0.35 | 0.04 |
| $J$=2–1 |  | 0.61 | 0.29 | 0.82 | 0.02 | R54/21 = 0.87
| $J$=5–4 | 1342213733 | 0.859 | 0.154 | 0.980 | 0.003 |
| $J$=9–8 | 1342216328 | 0.150 |  $\lesssim$ 0.006 | 0.293 | 0.003 | R98/54 = 0.24
| $J$=14–13 | 1342229786 |  $\lesssim$ 0.02 |  $\lesssim$ 0.02 |  $\lesssim$ 0.02 | 0.01 |
LSR velocity ranges (km s-1) |  | -7 – -1 | -19 – -13 | 13 – 19 |  |
NGC 6720 | $J$=1–0 |  | 0.18 | 0.30 |  $\lesssim$ 0.04 | 0.02 | R54/21 = 0.26
| $J$=2–1 |  | 0.36 | 1.57 |  $\lesssim$ 0.06 | 0.03 |
| $J$=5–4 | 1342210075 | 0.207 | 0.323 | 0.127 | 0.004 | R98/54 = 0.33
| $J$=9–8 | 1342220499 | 0.076 | 0.084 |  $\lesssim$ 0.007 | 0.003 |
LSR velocity ranges (km s-1) |  | 13 – 21 | -4.4 – 3.6 | 25.8 – 33.8 |  |
NGC 6781 | $J$=1–0 |  | 0.46 | 0.90 | 0.19 | 0.06 | R54/21 = 0.36
| $J$=2–1 |  | 1.43 | 2.60 |  $\lesssim$ 0.14 | 0.07 |
| $J$=5–4 | 1342230381 | 0.435 | 0.631 | 0.332 | 0.003 | R98/54 = 0.12
| $J$=9–8 | 1342229920 | 0.049 | 0.098 |  $\lesssim$ 0.006 | 0.003 |
LSR velocity ranges (km s-1) |  | -18.5 – -15.5 | -23.2 – -20.2 | -15 – -12 |  |
NGC 7293 | $J$=1–0 |  | 1.99 | 1.45 | 1.75 | 0.11 | R54/21 = 0.16
| $J$=2–1 |  | 1.65 | 0.97 | 1.07 | 0.01 |
| $J$=5–4 | 1342210082 | 0.265 | 0.131 | 0.251 | 0.004 | R98/54 $\lesssim$ 0.06
| $J$=9–8 | 1342220502 |  $\lesssim$ 0.007 |  $\lesssim$ 0.007 |  $\lesssim$ 0.007 | 0.004 |

###  4.2 Pseudo colour-colour diagrams

Figure 3: Pseudo colour-colour diagram comparing the $J$=5–4/$J$=2–1 and
$J$=9–8/$J$=5–4 intensity ratios of the wings for the compact sources and the
line profiles for the extended PNe. Although they are based on different line
ratios, some sources from Bujarrabal et al. (2012) are also included here (see
text for details). Figure 4: Pseudo colour-colour diagram comparing the
$J$=5–4/$J$=2–1 and $J$=9–8/$J$=5–4 intensity ratios of the wings and cores of
the compact sources (joined by arrows pointing towards the line cores) and the
identified velocity components of the extended PNe (connecting ratios of each
component to its average).

The relative contribution of the respective low-to-intermediate excitation CO
lines can be explored by mean of the corresponding line ratio, distinctively
formed for the various velocity intervals defined in the previous section. In
this study, we define R54/21 and R98/54 as the respective ratio of the $J$=5–4
/ $J$=2–1 and $J$=9–8 / $J$=5–4 characteristic intensities derived previously.
Their uncertainties were calculated using standard error propagation, assuming
absolute calibration errors of 20% for $J$=2–1, and 15% for $J$=5–4 and
$J$=9–8 (HIFI Handbook***https://www.cosmos.esa.int/web/herschel/legacy-
documentation-hifi). As explained in Section 4.1, the physical meaning of the
chosen velocity interval is different for the compact and extended sources in
our sample. This distinction also has implications in the analysis of the
ratios discussed here. For compact sources, we distinguish the ratios computed
over the line core intervals and those computed over the line wings intervals
– for the latter, we take an average of both the red and blue wings. For
extended sources, the ratios are computed for each separate velocity
component, as well as based on the averaged intensities over all considered
components. These respective ratios are compiled in Table 4 (compact sources)
and Table 5 (extended sources).

We then plot the R98/54 ratio against the R54/21 one, which we call in the
following pseudo colour-colour diagrams (thereafter cc-diagram). Although
different CO transitions were probed by Bujarrabal et al. (2012), we decided
to also represent in this diagram the ratios obtained for five sources (CRL
2688, CRL 618, IRAS 07134+1005, OH 231.8+4.2 and NGC 6302) by replacing the
$J$=5–4 line by the $J$=6–5 one, and the $J$=9–8 line by the $J$=10–9 one.
While not strictly comparable to our own ratios, we added those points to
check the general trends in the pseudo colour-colour space.

Figure 3 shows this diagram for the wings of the compact sources, and the
intensity averages of the extended objects. Three different groups emerge from
this plot: extended planetary nebulae, compact planetary nebulae and
protoplanetary nebulae. Extended planetary nebulae (in black) all show low
values of the average R98/54 ratio (typically below unity), while their R54/21
ratio spreads over an interval range from 0.16 to 0.9. Compact planetary
nebulae (in blue), however, seem to concentrate low values of R54/21 (all
below 0.2), while their R98/54 ratio is typically between 0.5 and 1, with the
exception of NGC 6537 (see discussion later in this paper). Protoplanetary
nebulae (in orange) also show a relatively low R54/21 ratio (around 0.1 on
average), but their R98/54 ratio, on the other hand, takes values that are
systematically in excess of unity.

Although they are not strictly based on the same metrics, the sources from
Bujarrabal et al. (2012) follow this general classification, but with a
certain deviation. On the one hand, CRL618 shows a noticeably higher R65/21
ratio (i.e. around 0.4), while OH231.8+4.2 only offers an upper limit to the
R109/65 ratio located well below 1. Still, the younger objects are mainly
located in the upper left corner of the diagram, while the planetary nebulae
tend to be in the lower area.

The overall trends observed in Fig. 3 suggest that R98/54 (and the equivalent
ratio formed from the $J$=10–9 over $J$=6–5 ratio for the sources in
Bujarrabal et al. 2012) is sensitive to the evolutionary stage of the nebulae.
This manifests as a higher excitation state in the high-velocity gas of pPNe
than in the shocked gas probed in the PNe. NGC 6537 stands as a noticeable
exception in that respect. Although considered a compact planetary nebula in
our sample, the optical and FIR maps (Fig. 39) indicates that the source is
partially resolved within the HIFI beam for the $J$=5–4 transition. Moreover,
if the different features we see are interpreted as separate velocity
components, their observational characteristics differ from those in the rest
of the extended object sample, as the spectra were collected very close to the
source photocentre (while most lines-of-sight in the extended source sample
pointed towards offset regions of the nebula). As such, NGC 6537 is not
directly comparable to either of the compact or extended PN source samples and
this fact could explain its outlying location on the diagram.

We also note that the wings of 89 Her, R Sct, OH 17.7–2.0 and M 2–56 are not
detected at high-$J$ (they are not featured in Fig.3, but their ratios in the
line core are shown in Fig. 4). As we mentioned in Sect. 3, 89 Her and R Sct
are considered bipolar nebulae with a rotational disk. It is believed that
this type of nebula concentrates all its gas in the disk and central region
and, as such, they are considered compact objects suffering from severe
dilution in the beam of the telescope.

A comparison of the line ratios as measured in the line wings and line core of
the compact objects is shown in Fig. 4. Each pair of ratios is joined with an
arrow pointing toward the line cores, allowing us to trace possible trends
between the relative excitation of the outflows with regard to that of the
cores. We note that the arrows of the pPNe sub-sample have greater amplitude
and point downwards, indicating that the temperature of their fast outflows
must be higher than in the slow winds and suggesting, therefore, a more recent
gas acceleration. On the contrary, the difference between the ratios in the
line wings and the line cores of the compact PNe sub-sample is smaller and
most of their arrows point upwards, indicative of cooler and older gas in the
fast flows. This trend is similar to that reported in Bujarrabal et al. (2012)
for the objects of their sample represented in our pseudo colour-colour
diagram.

Figure 4 also includes a visualisation of the distribution of the line ratios
for the various components of the extended PNe sub-samples. In this case, the
ratio in each component is connected to that of the averaged ratios shown in
Fig. 3. At most, this should indicate whether the individual velocity
components in each object do or do not share certain conditions of excitation
or evolution. Overall, there is no clear trend derived from them. In the case
of NGC 6781, all three components seem to behave similarly, while NGC 2346,
for example, exhibits a larger spread of physical conditions in the line
ratios.

###  4.3 Comparison with theoretical predictions

We compare in this section the distribution of our line ratios with
theoretical predictions in order to estimate the physical conditions in the
fast outflows of our source sample. To this aim, we have performed simplified
calculation based on the large velocity gradient (LVG) approximation, for a
grid of density, $n_{\rm tot}$, and kinetic temperature, $T_{\rm k}$, as well
as different optical depth regimes. This assumption is justified by the
kinematics at play in pPNe and many PNe. We also assume that no strong
radiative interactions take place between distant points of the envelope
exhibiting very different physical conditions. The code is similar to the one
used by Bujarrabal et al. (2012); we have updated the calculation of the
collisional transition probabilities, including, in particular, a more
accurate interpolation between rates for different temperatures and allowing
for calculations at a higher number of levels.

Figure 5: Pseudo colour-colour diagram comparing the $J$=5–4/$J$=2–1 and
$J$=9–8/$J$=5–4 intensity ratios of pPNe line wings, compact PNe line wings
and the LVG models for different optical depths (solid, dotted and dashed
lines), densities (number indicates log10($n_{\rm tot}$)) and temperatures
(coded as shown in the colour scale). See text for further details.

Our computations are based on canonical envelope properties assuming the
following parameters: a distance to the central star of $R$ = 6$\times$1016
cm, an expansion velocity of $V_{\rm exp}$ = 15 km s-1 and a relative
abundance of $X(^{12}$CO) = 2$\times$ 10-4 (for more details, see Sect. 3.1.
of Bujarrabal et al. 2012). Opacity effects are taken into account in the
calculations under the usual assumptions of the LVG approximation. Although
such an assumption implies a constant value for $R$ $X/V_{\rm exp}$ (where
$n_{tot}R$ $X/V_{\rm exp}$ represents the column density opacity under the LVG
approximation), this approach offers an empirical set of models which is
sufficient to probe the statistical behaviour of our line ratios in this grid.
Similarly, because we want to locate all our objects in the same set of model
space, we do not include any ad hoc modelling of the large-scale structure of
the nebula.

We also assume that the lines come from a region significantly smaller than
the telescope spatial resolution, which applies to our compact source sample.
For comparison with our extended nebulae, we corrected our line ratio
calculations, taking into account the ratios of the applicable telescope beams
(Table 2).

Figure 5 shows our predicted line ratios for several optical depths,
densities, and temperature regimes within the ranges relevant to our data. The
theoretical results for different temperatures are presented using a colour
coding where the adopted scale ranges from 10 K to 200 K. As anticipated in
Section 4.2, the 12CO $J$=9–8 / $J$=5–4 ratio appears as a good proxy for the
kinetic temperature, especially in the optically thick regime. The slope of
this relationship is also related to the density. Conversely, the 12CO $J$=5–4
/ $J$=2–1 ratio appears to be inversely proportional to both optical depth and
density. Unfortunately, predictions from the various model grids overlap in
the pseudo colour-colour diagram, leading to some degeneracy in the derivation
of the physical conditions from our measurements. Figure 5 also shows the
intensity ratios of the line wings for the compact sources.

From the comparison of our pPN observational results with the LVG predictions,
we observe that high density optically thick models predict too low
temperatures (¡ 25 K) for the outflows of this kind of objects, while the
optically thin approximations assign densities excessively high (larger than
107 cm-3). Consequently, we used the model grid based on an intermediate
optical depth and optically thick models of 104 cm-3, for which the fast
outflows of protoplanetary nebulae are compatible with densities between 104
and 106 cm-3 and temperatures in the range of 75–200 K. Doing the same
exercise for the compact PNe, we derive lower kinetic temperatures in the
25–75 K range and densities in the 104–106 cm-3 range.

We interpret this result as a relation between the temperature of the high-
velocity gas found in the nebulae and the time elapsed since its acceleration
due to shocks. The heated and accelerated gas reaches high temperatures during
the ejection of very fast outflows, but it cools in a short period of time.
This is consistent with the results already reported by Bujarrabal et al.
(2012). The correlation between the density of the warm regions and the
evolutionary stage of the nebulae, however, is not so clear.

###  4.4 Intermediate velocity ranges

Aside from the main spectral components analysed above, we note the existence
of intermediate velocity features in some of our sources, particularly in M
1–92 and M 1–17. Following the same methods described in the previous
sections, we derived the characteristic intensities and respective $R_{54/21}$
and $R_{98/54}$ ratios of these intermediate line regions for the two
mentioned nebulae. Their values are collected in Table 6, alongside the LSR
velocity ranges adopted (also marked in Figs. 9 and 16). The substantial
difference between the ratios of these regions and those estimated in the line
wings is indicative of a cooler gas component of intermediate velocity in the
nebulae that, in the case of M 1–92, could correspond with the lobe walls
identified in Alcolea et al. (2007) and Bujarrabal et al. (1998).

Table 6: Characteristic intensities of two intermediate velocity ranges
identified in the CO profiles of the pPNe M 1–17 and M 1–92, and their derived
ratios

Source |  12CO line |  | Blue interm. wing
---
av. Tmb(K)
| Red interm. wing
---
av. Tmb(K)
| Noise
---
r.m.s. (K)
| Line ratios
---
Blue interm. line wings
| Line ratios
---
Red interm. line wings
LSR velocity ranges (km s-1) | -25 – -10.5 | 10.5 – 25 |  |  |
M 1–92 | $J$=1–0 | 0.36 | 0.43 | 0.01 | R54/21 = 0.05 | R54/21 = 0.04
| $J$=2–1 | 0.88 | 1.18 | 0.02 |  |
| $J$=5–4 | 0.048 | 0.052 | 0.002 | R98/54 = 0.55 | R98/54 = 0.68
| $J$=9–8 | 0.026 | 0.035 | 0.002 |  |
LSR velocity ranges (km s-1) | 11.5 – 23.5 | 32 – 44 |  |  |
M 1–17 | $J$=1–0 | 1.06 | 1.11 | 0.03 | R54/21 = 0.06 | R54/21 = 0.08
| $J$=2–1 | 1.86 | 1.89 | 0.04 |  |
| $J$=5–4 | 0.112 | 0.144 | 0.005 | R98/54 = 1.19 | R98/54 = 0.94
| $J$=9–8 | 0.133 | 0.135 | 0.004 |  |

##  5 Conclusions

We used the high-resolution spectrometer HIFI on board Herschel to survey a
sample of nine protoplanetary nebulae (pPNe) and nine planetary nebulae (PNe)
in intermediate-excitation molecular lines of 12CO ($J$=5–4, $J$=9–8, and
$J$=14–13 for a subset of those). We complemented these data with observations
of the low-$J$ transitions of CO using IRAM 30m telescope and literature
(Table 1). The selected lines are shown in Figs. 6 to 21 and described briefly
in Sect. 3.

As in Bujarrabal et al. (2012), we carried out a simplified analysis of the CO
lines in terms of the emission of three representative profile components. The
physical significance behind these regions depends on the relative size of the
objects with regard to the dimension of HIFI beams. Five of the PNe in our
sample were considered extended sources and we interpreted their three main
profile features as three different velocity components. On the other hand,
the representative profile features of the compact nebulae were analysed as a
line core (known to come from the remnant AGB shell and rotating disks in two
cases) and two line wings (associated with the fast outflows).

In Tables 4 and 5, we compiled the averaged intensities of the characteristic
components found in compact and extended sources respectively, together with
their errors. These tables indicate the intensity ratios $J$=5–4/$J$=2–1 and
$J$=9–8/$J$=5–4 derived for the line core and line wings of compact objects,
and for the entire line profile of extended nebulae. These ratios are
represented in Figs. 3 and 4 as so-called pseudo colour-colour diagrams. In
these diagrams, our sample is clearly divided with regard to their
evolutionary stage and to their relative extension. While the wings of pPNe
are located at the top left of the diagram and are significantly more excited
than their cores, the wings of PNe are found at its bottom left and are less
excited than their cores. This observation suggests a correlation between the
evolutionary stage of nebulae and the excitation state of their fast outflows.
With respect to the apparent spatial extension, compact objects remain at the
left of the pseudo cc-diagram, while the extended sources are found at its
bottom right.

We performed radiative transfer calculations based on the LVG approximation to
derive the physical properties of the nebular gas. We found that the kinetic
temperatures of the fast gas presented in pPNe tend to be higher (with typical
values between 75-200 K) than those found in PNe (with characteristic
temperatures between 25 and 75 K). This result indicates a relation between
the temperature of fast outflows and the time elapsed since their ejection. It
is interpreted as the effect of cooling in the heated and accelerated gas
during the pPN phase, when the shocks propagate across the nebula. This
conclusion is strengthened by the significantly high excitation state found
for the high-velocity line wings in M 1–92 (much higher than for the
intermediate-velocity wings); the high-velocity emission is indeed known to
come from the tips of the nebula lobes, which are directly impinged by the
post-AGB jets.

###### Acknowledgements.

We thank the anonymous referee for their useful and constructive comments on
the paper. We also thank Prof. Hans Olofsson for his thorough review of this
manuscript. HIFI has been designed and built by a consortium of institutes and
university departments from across Europe, Canada and the USA under the
leadership of SRON Netherlands Institute for Space Research, Groningen, The
Netherlands, and with major contributions from Germany, France and the USA.
Consortium members are: Canada: CSA, U.Waterloo; France: CESR, LAB, LERMA,
IRAM; Germany: KOSMA, MPIfR, MPS; Ireland: NUI Maynooth; Italy: ASI, IFSI-
INAF, Osservatorio Astrofisico di Arcetri-INAF; The Netherlands: SRON, TUD;
Poland: CAMK, CBK; Spain: Observatorio Astronómico Nacional (IGN), Centro de
Astrobiología (CSIC-INTA); Sweden: Chalmers University of Technology – MC2,
RSS & GARD; Onsala Space Observatory; Swedish National Space Board, Stockholm
University – Stockholm Observatory; Switzerland: ETH Zurich, FHNW; USA:
Caltech, JPL, NHSC. This paper is based on obser-vations carried out at the
IRAM-30 m single-dish telescope. IRAM is sup-ported by INSU/CNRS (France), MPG
(Germany) and IGN (Spain). JA and VB are partly funded by the Spanish MICIIN
grant AYA2016-78994-P.

## References

  * Alcolea & Bujarrabal (1995) Alcolea, J. & Bujarrabal, V. 1995, A&A, 303, L21
  * Alcolea et al. (2001) Alcolea, J., Bujarrabal, V., Sánchez Contreras, C., Neri, R., & Zweigle, J. 2001, A&A, 373, 932
  * Alcolea et al. (2007) Alcolea, J., Neri, R., & Bujarrabal, V. 2007, A&A, 468, L41
  * Ashley & Hyland (1988) Ashley, M. C. B. & Hyland, A. R. 1988, ApJ, 331, 532
  * Bachiller et al. (1989) Bachiller, R., Bujarrabal, V., Martin-Pintado, J., & Gomez-Gonzalez, J. 1989, A&A, 218, 252
  * Bachiller et al. (1997) Bachiller, R., Forveille, T., Huggins, P. J., & Cox, P. 1997, A&A, 324, 1123
  * Bachiller et al. (1991) Bachiller, R., Huggins, P. J., Cox, P., & Forveille, T. 1991, A&A, 247, 525
  * Bachiller et al. (1993) Bachiller, R., Huggins, P. J., Cox, P., & Forveille, T. 1993, A&A, 267, 177
  * Bains et al. (2003) Bains, I., Gledhill, T. M., Yates, J. A., & Richards, A. M. S. 2003, MNRAS, 338, 287
  * Balick & Frank (2002) Balick, B. & Frank, A. 2002, ARA&A, 40, 439
  * Bublitz et al. (2019) Bublitz, J., Kastner, J. H., Santander-García, M., et al. 2019, A&A, 625, A101
  * Bujarrabal et al. (1998) Bujarrabal, V., Alcolea, J., & Neri, R. 1998, ApJ, 504, 915
  * Bujarrabal et al. (1992) Bujarrabal, V., Alcolea, J., & Planesas, P. 1992, A&A, 257, 701
  * Bujarrabal et al. (2012) Bujarrabal, V., Alcolea, J., Soria-Ruiz, R., et al. 2012, A&A, 537, A8
  * Bujarrabal et al. (2013) Bujarrabal, V., Alcolea, J., Van Winckel, H., Santand er-García, M., & Castro-Carrizo, A. 2013, A&A, 557, A104
  * Bujarrabal et al. (2001) Bujarrabal, V., Castro-Carrizo, A., Alcolea, J., & Sánchez Contreras, C. 2001, A&A, 377, 868
  * Bujarrabal et al. (1994) Bujarrabal, V., Fuente, A., & Omont, A. 1994, A&A, 285, 247
  * Bujarrabal et al. (2007) Bujarrabal, V., van Winckel, H., Neri, R., et al. 2007, A&A, 468, L45
  * Calvet & Cohen (1978) Calvet, N. & Cohen, M. 1978, MNRAS, 182, 687
  * Castro-Carrizo et al. (2002) Castro-Carrizo, A., Bujarrabal, V., Sánchez Contreras, C., Alcolea, J., & Neri, R. 2002, A&A, 386, 633
  * Castro-Carrizo et al. (2010) Castro-Carrizo, A., Quintana-Lacaci, G., Neri, R., et al. 2010, A&A, 523, A59
  * Cox et al. (1992) Cox, P., Omont, A., Huggins, P. J., Bachiller, R., & Forveille, T. 1992, A&A, 266, 420
  * Cuesta et al. (1995) Cuesta, L., Phillips, J. P., & Mampaso, A. 1995, A&A, 304, 475
  * Danilovich et al. (2015) Danilovich, T., Teyssier, D., Justtanont, K., et al. 2015, A&A, 581, A60
  * De Beck et al. (2010) De Beck, E., Decin, L., de Koter, A., et al. 2010, A&A, 523, A18
  * de Graauw et al. (2010) de Graauw, T., Helmich, F. P., Phillips, T. G., et al. 2010, A&A, 518, L6
  * de Ruyter et al. (2006) de Ruyter, S., van Winckel, H., Maas, T., et al. 2006, A&A, 448, 641
  * Edwards et al. (2014) Edwards, J. L., Cox, E. G., & Ziurys, L. M. 2014, ApJ, 791, 79
  * Edwards & Ziurys (2013) Edwards, J. L. & Ziurys, L. M. 2013, ApJ, 770, L5
  * Fong et al. (2006) Fong, D., Meixner, M., Sutton, E. C., Zalucha, A., & Welch, W. J. 2006, ApJ, 652, 1626
  * Frew (2008) Frew, D. J. 2008, PhD thesis, Department of Physics, Macquarie University, NSW 2109, Australia
  * Gledhill (2005) Gledhill, T. M. 2005, MNRAS, 356, 883
  * Gledhill et al. (2001) Gledhill, T. M., Chrysostomou, A., Hough, J. H., & Yates, J. A. 2001, MNRAS, 322, 321
  * Heske et al. (1990) Heske, A., Forveille, T., Omont, A., van der Veen, W. E. C. J., & Habing, H. J. 1990, A&A, 239, 173
  * Hrivnak et al. (1989) Hrivnak, B. J., Kwok, S., & Volk, K. M. 1989, ApJ, 346, 265
  * Huggins et al. (2000) Huggins, P. J., Forveille, T., Bachiller, R., & Cox, P. 2000, ApJ, 544, 889
  * Justtanont et al. (2000) Justtanont, K., Barlow, M. J., Tielens, A. G. G. M., et al. 2000, A&A, 360, 1117
  * Kwok et al. (2010) Kwok, S., Chong, S.-N., Hsia, C.-H., Zhang, Y., & Koning, N. 2010, ApJ, 708, 93
  * Kwok et al. (2008) Kwok, S., Chong, S.-N., Koning, N., Hua, T., & Yan, C.-H. 2008, ApJ, 689, 219
  * Kwok et al. (1989) Kwok, S., Volk, K. M., & Hrivnak, B. J. 1989, ApJ, 345, L51
  * Le Bertre (1987) Le Bertre, T. 1987, A&A, 180, 160
  * Maciel (1984) Maciel, W. J. 1984, A&AS, 55, 253
  * Matsuura et al. (2005) Matsuura, M., Zijlstra, A. A., Gray, M. D., Molster, F. J., & Waters, L. B. F. M. 2005, MNRAS, 363, 628
  * Meaburn et al. (2005) Meaburn, J., Boumis, P., López, J. A., et al. 2005, MNRAS, 360, 963
  * Meixner et al. (1997) Meixner, M., Skinner, C. J., Graham, J. R., et al. 1997, Astronomy Data Image Library
  * Milam et al. (2009) Milam, S. N., Woolf, N. J., & Ziurys, L. M. 2009, ApJ, 690, 837
  * O’Dell et al. (2013) O’Dell, C. R., Ferland, G. J., Henney, W. J., & Peimbert, M. 2013, AJ, 145, 170
  * Omont et al. (1995) Omont, A., Moseley, S. H., Cox, P., et al. 1995, ApJ, 454, 819
  * Sánchez Contreras et al. (2004) Sánchez Contreras, C., Bujarrabal, V., Castro-Carrizo, A., Alcolea, J., & Sargent, A. 2004, ApJ, 617, 1142
  * Sánchez Contreras et al. (2007) Sánchez Contreras, C., Le Mignant, D., Sahai, R., Gil de Paz, A., & Morris, M. 2007, ApJ, 656, 1150
  * Schmidt et al. (2018) Schmidt, D. R., Zack, L. N., & Ziurys, L. M. 2018, ApJ, 864, L31
  * Schmidt & Ziurys (2016) Schmidt, D. R. & Ziurys, L. M. 2016, ApJ, 817, 175
  * Shenton et al. (1995) Shenton, M., Evans, A., & Williams, P. M. 1995, MNRAS, 273, 906
  * Shibata et al. (1989) Shibata, K. M., Tamura, S., Deguchi, S., et al. 1989, ApJ, 345, L55
  * Su et al. (2004) Su, K. Y. L., Kelly, D. M., Latter, W. B., et al. 2004, ApJS, 154, 302
  * Teyssier et al. (2011) Teyssier, D., Alcolea, J., Bujarrabal, V., et al. 2011, 280, 353
  * Ueta et al. (2000) Ueta, T., Meixner, M., & Bobrowsky, M. 2000, ApJ, 528, 861
  * Van Winckel & Reyniers (2000) Van Winckel, H. & Reyniers, M. 2000, A&A, 354, 135
  * Waters et al. (1993) Waters, L. B. F. M., Waelkens, C., Mayor, M., & Trams, N. R. 1993, A&A, 269, 242
  * Young et al. (1999) Young, K., Cox, P., Huggins, P. J., Forveille, T., & Bachiller, R. 1999, ApJ, 522, 387
  * Zack & Ziurys (2013a) Zack, L. N. & Ziurys, L. M. 2013a, ApJ, 765, 112
  * Zack & Ziurys (2013b) Zack, L. N. & Ziurys, L. M. 2013b, ApJ, 765, 112
  * Zeigler et al. (2013) Zeigler, N. R., Zack, L. N., Woolf, N. J., & Ziurys, L. M. 2013, ApJ, 778, 16
  * Zhang et al. (2008) Zhang, Y., Kwok, S., & Dinh-V-Trung. 2008, ApJ, 678, 328
  * Ziurys et al. (2020) Ziurys, L. M., Schmidt, D. R., & Woolf, N. J. 2020, ApJ, 900, L31

##  Appendix A CO spectra of the source sample

In this section, we compile all the CO transition spectra used in this work,
together with the adopted systemic velocity and velocity ranges considered in
the analysis. See the captions in Figures 1 and 2 for further details.

Figure 6: Observations of IRAS 07134+1005. Shadowed regions mark the LSR
velocity ranges adopted for each feature and the dashed vertical line
indicates the systemic velocity. Figure 7: Observations of IRAS 19500–1709.
Shadowed regions mark the LSR velocity ranges adopted for each feature and the
dashed vertical line indicates the systemic velocity. Figure 8: Observations
of IRAS 22272+5435. Shadowed regions mark the LSR velocity ranges adopted for
each feature and the dashed vertical line indicates the systemic velocity.
Figure 9: Observations of M 1–92. Shadowed regions mark the LSR velocity
ranges adopted for each feature and the dashed vertical line indicates the
systemic velocity. See Section 4.4 for the reference to the intermediate
wings. Figure 10: Observations of M 2–56. Shadowed regions mark the LSR
velocity ranges adopted for each feature and the dashed vertical line
indicates the systemic velocity. Figure 11: Observations of OH 17.7–2.0.
Shadowed regions mark the LSR velocity ranges adopted for each feature and the
dashed vertical line indicates the systemic velocity. Figure 12: Observations
of R Sct. Shadowed regions mark the LSR velocity ranges adopted for each
feature and the dashed vertical line indicates the systemic velocity. The CO
$J$=3–2 line is taken from De Beck et al. (2010). It is shown here only as an
illustration of the line profile evolution between the low- and
intermediate-$J$ CO lines, but it will not be used in our quantitative study.
Figure 13: Observations of 89 Her. Shadowed regions mark the LSR velocity
ranges adopted for each feature and the dashed vertical line indicates the
systemic velocity. Figure 14: Observations of NGC 6537. Shadowed regions mark
the LSR velocity ranges adopted for each feature and the dashed vertical line
indicates the systemic velocity. The CO $J$=3–2, $J$=4–3 and $J$=6–5
transitions are taken from Edwards et al. (2014). They are shown here just as
an illustration of the line profile evolution between the low- and
intermediate-$J$ CO lines, but they are used in our quantitative study. Figure
15: Observations of M 1–16. Shadowed regions mark the LSR velocity ranges
adopted for each feature and the dashed vertical line indicates the systemic
velocity. Figure 16:  Observations of M 1–17. Shadowed regions mark the LSR
velocity ranges adopted for each feature and the dashed vertical line
indicates the systemic velocity. See Section 4.4 for the reference to the
intermediate wings. Figure 17: Observations of M 3–28. Shadowed regions mark
the LSR velocity ranges adopted for each feature and the dashed vertical line
indicates the systemic velocity. The CO $J$=3–2, transitions is taken from
Schmidt & Ziurys (2016). It is shown here just as an illustration of the line
profile evolution between the low- and intermediate-$J$ CO lines, but is not
used in our quantitative study. Figure 18: Observations of NGC 2346. Shadowed
regions indicate the velocity ranges adopted for the detected velocity
components and the dashed line marks the systemic velocity. Figure 19:
Observations of NGC 6072. Shadowed regions indicate the velocity ranges
adopted for the detected velocity components and the dashed line marks the
systemic velocity. Figure 20: Observations of NGC 6781. Shadowed regions
indicate the velocity ranges adopted for the detected velocity components and
the dashed line marks the systemic velocity. Figure 21: Observations of NGC
7293. Shadowed regions indicate the velocity ranges adopted for the detected
velocity components and the dashed line marks the systemic velocity.

##  Appendix B Optical and FIR images of the sample

In this section, we present the optical and FIR images (when available) of our
targets. The Herschel/HIFI beams of the lines used in this work are overlayed
on these maps red contours, in order of size from largest to smallest: 12CO
$J$=5–4, 12CO $J$=9–8 and 12CO $J$=14–13.

Figure 22: Image of IRAS 07134+1005 obtained from DSS2 colour.

Figure 23: FIR image of IRAS 07134+1005 obtained by Herschel PACS RGB 70,160
microns.

Figure 24: Image of IRAS 19500–1709 obtained from DSS2 colour.

Figure 25: Image of IRAS 21282+5050 obtained from DSS2 colour.

Figure 26: Image of IRAS 22272+5435 obtained from DSS2 colour.

Figure 27: FIR image of IRAS 22272+5435 obtained by Herschel PACS RGB 70,160
microns.

Figure 28: Image of M 1–92 obtained from DSS2 colour.

Figure 29: FIR image of M 1–92 obtained by Herschel PACS RGB 70,160 microns.

Figure 30: Image of M 2–56 obtained from DSS2 colour.

Figure 31: Image of OH 17.7–2.0 obtained from DSS2 colour.

Figure 32: Image of R Sct obtained from DSS2 colour.

Figure 33: Image of 89 Her obtained from DSS2 colour.

Figure 34: NASA/ESA Hubble Space Telescope image of NGC 6537, a composite of
[N II] 6584 $\AA$ (red), H$\alpha$ 6563 $\AA$ (green), and [O III] 5007 $\AA$
(blue) emission.

Figure 35: FIR image of NGC 6537 obtained by Herschel PACS RGB 70,160 microns.

Figure 36: Image of M 1–16 obtained from DSS2 colour.

Figure 37: Image of M 1–17 obtained from DSS2 colour.

Figure 38: Image of M 3–28 obtained from DSS2 colour.

Figure 39: FIR image of M 3–28 obtained by Herschel PACS RGB 70,160 microns.

Figure 40: Composite image of NGC 6072, H$\alpha$ (red), [OIII] (green), B
filters (blue), produced by Ganymed Telescope.

Figure 41: NASA/ESA Hubble Space Telescope image of NGC 2346, a composite of R
band 6750 $\AA$ (red), H$\alpha$ 6560 $\AA$ (green), [SII] 6730 $\AA$ (blue).

Figure 42: NASA/ESA Hubble Space Telescope image of NGC 6720, a composite of R
filters 6580 $\AA$ and 6730 $\AA$ (red), G filters 6560 $\AA$ and 6450 $\AA$
and V filters 5020 $\AA$ (green), B filters 4690 $\AA$ and 4870 $\AA$ (blue).

Figure 43: FIR image of NGC 6720 obtained by Herschel PACS RGB 70,160 microns.

Figure 44: ESO image of NGC 6781, a composite of R and H$\alpha$ (red), V band
and [OIII] (green), B band (blue).

Figure 45: FIR image of NGC 6781 obtained by Herschel PACS RGB 70,160 microns.

Figure 46: Image of NGC 7293 obtained from DSS2 colour.

Figure 47: FIR image of NGC 7293 obtained by Herschel PACS RGB 70,160 microns.