\documentclass[journal]{rmaa}

\usepackage{multirow}
\newcommand{\casitiny}{\fontsize{6}{7}\selectfont}


\title{Echelle spectroscopy of the nuclei of the highly collimated bipolar planetary nebulae M~2-9 and M~1-91\altaffilmark{1}} 

\author{S. Torres-Peimbert,\altaffilmark{2}
A. Arrieta,\altaffilmark{3} 
and M. Bautista\altaffilmark{4}}

\altaffiltext{1}{Based on observations obtained at the Observatorio Astron\'omico Nacional at 
San Pedro M\'artir, Baja California, Mexico.}
\altaffiltext{2}{Instituto de Astronom\'ia, Universidad Nacional Aut\'o\-noma de M\'exico, Mexico.}
\altaffiltext{3}{Depto. de F\'isica y Matem\'aticas, Universidad Iberoamericana, Mexico.}
\altaffiltext{4}{Department of Physics, Western Michigan University, USA.}


\shorttitle{The planetary nebulae M~2-9 and M~1-91}

\shortauthor{Torres-Peimbert, Arrieta, \& Bautista}


\fulladdresses{
\item A. Arrieta:  Universidad Iberoamericana, Depto. de F\'isica y Matem\'aticas, Prol. Paseo de la Reforma 880, Lomas de Santa Fe, 
M\'exico D. F. 01219, Mexico (anabel.arrieta@uia.mx).
\item M. Bautista: Department of Physics, Western Michigan University, Kalamazoo, MI 49008, USA.
\item S. Torres-Peimbert: Instituto de Astronom\'{\i}a, Universidad Nacional Aut\'onoma de M\'exico, Apdo. Postal 70-264, M\'exico, D. F. 
04510, Mexico (silvia@astroscu.unam.mx).}


\listofauthors{S.~Torres-Peimbert, A.~Arrieta, \& M.~Bautista}

\indexauthor{Torres-Peimbert, S.}
\indexauthor{Arrieta, A.}
\indexauthor{Bautista, M.}


\ReceivedDate{2010 January 21} 
\AcceptedDate{2010 May 20} 
\SetYear{2010}
\SetVolume{46}


\resumen{Presentamos espectroscop\'ia echelle del n\'ucleo sin resolver de las nebulosas 
bipolares M~2-9 and M~1-91. Los espectros est\'an dominados por l\'ineas de emisi\'on 
emitidas en un amplio intervalo de condiciones f\'isicas. De las observaciones identificamos 
las l\'ineas de emisi\'on, las condiciones f\'isicas y los movimientos relativos de las 
diferentes especies ionizadas en la regi\'on circunestelar de ambos objetos. Proponemos 
que las l\'ineas prohibidas observadas se originan en la parte interior del toro extendido 
que rodea a cada objeto.}


\abstract{We present echelle spectroscopy of the unresolved nuclei of the bipolar planetary 
nebulae M~2-9 and M~1-91. The spectrum is dominated by emission lines emitted under a wide 
range of physical conditions. From the observations we identify the emission lines, derive 
physical conditions and relative motions of the different ionized species in the circumstellar 
region of both objects. We propose that the observed forbidden lines arise in the inner part 
of the extended torus that surrounds each object.}


\addkeyword{planetary nebulae: individual (M~2-9, M~1-91)}


\begin{document}

\maketitle


\section{Introduction}

Among the many interesting proto-planetary nebulae and young planetary nebulae there are 
those with extreme bipolar shape, or ``butterfly'' type. {M~2-9} and {M~1-91} are analogous 
not only in morphology, but also in dusty lobes and spectral characteristics of their nuclear 
regions. {M~2-9} has been widely studied in many of its characteristics, and yet is not fully 
understood; the studies on {M~1-91}, have been more limited because it is considerably fainter, 
and thus it is less well understood. We present here the result of our spectroscopic 
observations of their central objects in an attempt to further understand the physical 
conditions of the circumstellar gas. In what follows we will present each of the objects, 
describe the observations and their characteristics. For {M~2-9} we discuss the profiles of 
the Balmer lines, of [O III] lines, and the intensities of the Fe lines. We also derive the 
reddening law and the behavior of its physical conditions as a function of velocity relative to 
the central star. For {M~1-91} we carry out a similar study; however, as expected, the information 
derived for this object is more limited.


\section{Observations} 
\label{sobserv}

The observations were carried out at the Observatorio Astr\'onomico Nacional in San Pedro 
M\'artir, Baja California with the 2.1 m telescope and the REOSC echelle spectrograph 
(R $\sim$ 18,000 at 5,000~\AA) and a $1024 \times 1024$ Tektronix detector that yields a 
spectral resolution of 10.6~km~s$^{-1}$ and a spatial resolution of 0.99\arcsec\ per pixel. The 
3600 to 6850~\AA\ range was covered in 29 orders (blue spectrum) for both objects, and the 
6500 to 9500~\AA\ range in 11 orders (red spectrum) for {M~2-9}. 

Long slit observations, $79.8\arcsec \times 2\arcsec$, oriented along the major axis were 
obtained for both objects. The exposure time for the blue spectrum of {M~2-9} was 600~sec 
and the other spectra required 900 sec. 

Short window extractions (10 pixels) were performed to obtain the nuclear spectra, which will 
be reported here. For {M~2-9} the blue spectrum was observed on 23 April 1995 and the red on 
25 April 1995; while for {M~1-91} the observations were carried out on 27 May 1998. The data 
reduction was carried out with IRAF routines. The spectra were normalized to the continuum for 
each spectral order with a Chebyshev polynomial.


\section{M~2-9}

One of the most interesting objects is {M~2-9} (PN G10.8+18), a bright nebula of extreme bipolar 
morphology. It is seen almost face-on, with a size of approximately $12\arcsec \times 48\arcsec$ and extensions  
that reach up to 60\arcsec\ from the nucleus (Schwarz et al. 1997). It has been extensively 
studied (cf. Allen \& Swings 1972; Balick 1989; Hora \& Latter 1994; Schwarz et al. 1997; Solf 
2000; Doyle et al. 2000). 

The evolutionary status of {M~2-9} is still controversial. {M~2-9} has been identified with objects 
which ``bear the characteristics indicative for very young planetary nebulae'', such as {M~1-91}, 
M~1-92 (Solf 1984), CRL 618 (Schmidt \& Cohen 1981) or ``proto-planetary nebulae'' (Walsh 1981). 
Schwarz et al. (1997) derive a distance to {M~2-9} of 650 pc. At this distance the bolometric 
luminosity is 550~$L_\odot$ and its dynamic age is 1200 years. So at least one Mira-containing 
symbiotic system is reminiscent of the odd behavior pattern of M~2-9 (Doyle et al. 2000). 

The lobes are found to have a double shell structure; the outer shell is dominated by H$_2$ line 
emission (Hora \& Latter 1994) The lobes have condensations symmetrically located relative to the 
equatorial plane. The weak emissions observed at 60\arcsec\ S and N of the nucleus are also located 
in a point-symmetric distribution. One of the more significant characteristics of this object is 
that the knots brighter and closer to the nucleus show motions parallel to the equatorial plane 
(Allen \& Swings 1972; Kohoutek \& Surdej 1980; Balick 1989; Doyle et al. 2000).


%%%%% FIGURA 1

\begin{figure}[!t]
\vspace{0.1cm}
\centering
\includegraphics[width=0.99\columnwidth]{storresfig1.eps}
\caption{HST image of the bipolar planetary nebula {M~2-9} from Balick et al. (1997). 
Also shown are details of the internal region taken with two different filters: [OIII] on the 
left, and [S~II] on the right.}
\label{fm29image}
\vspace{0.3cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Polarimetric studies show that the radiation from the lobes is polarized, suggesting that the nebular 
emission has two components: one produced locally in the lobes and another one probably being light 
scattered from the bright central region (Trammell, Goodrich, \& Dinerstein 1995). 

Upon examination of direct HST images Balick (1999) found that the [O III] image shows a series of 
knots lying along the nebula's symmetry axis. Unlike other knots in the lobes, the near-nuclear [O III] 
knots have no counterparts in other emission lines. The enlarged [S II] image shows a network of nearly 
radial filaments that are not artifacts of the telescope optics. These are associated with the base 
of the lobes. Figure~\ref{fm29image} shows an image of the bipolar nebula in the light of [O~III] and 
[S~II]. An enlargement of the central regions is presented (from Balick, Icke, \& Mellema 1997).


\subsection{Nucleus of M~2-9}
\label{sm29}

{M~2-9} is a very dust rich nebula. It has been found that the reddening of the nucleus is significantly 
different from that shown in the lobes. In the nucleus, different authors have found different reddening 
values; Swings \& Andrillat (1979) and Calvet \& Cohen (1978) determined an $A_\mathrm{v} \geq 5$ mag, while 
Hora \& Latter (1994) from the hydrogen Brackett series determined an $A_\mathrm{v}=2.3$ mag.

The nuclear spectrum shows nebular lines in emission over a weak continuum. No absorption lines have 
been observed, therefore the determination of its spectral type is very uncertain. It has been proposed 
that the nucleus is composed of a binary system: a hot component that ionizes the O$^{++}$ gas and a 
B5 star that contributes to its visible luminosity (Calvet \& Cohen 1978). Also Schwarz et al. (1997) 
suggest that a compact but hot component must be present; the hot component ionizes the high 
excitation emission lines while the B5 (or cooler) star provides part of the luminosity and the binary 
interaction of the pair produces the high collimation of the nebula. An accretion disk around the compact 
star would provide the fast disk wind.

Among other possible explanations, Livio \& Soker (2001) have proposed a model to explain the apparent 
motion of the knots and the evolutionary stage of the exciting star. They propose that the nucleus is 
constituted by a binary system composed by an AGB or post-AGB star and a white dwarf.

Kwok et al. (1985) observed at five different radio continuum frequencies angular radii of the optically 
thick surface from 0.033\arcsec\ to 22\arcsec, depending on the frequency. Adopting a value of the wind 
velocity of 1600~km~s$^{-1}$ they derive a mass loss rate of $3.3 \times 10^{-5}\, M_{\odot}$~yr$^{-1}$ for a 
distance of 1 kpc. If the distance is 650 pc, then the mass loss rate is $1.1 \times 10^{-5}\, M_\odot$~yr$^{-1}$. 
This mass loss rate is much higher than the rates ($< 10^{-7} \,M_\odot$~yr$^{-1}$) estimated for other 
planetary nebulae. They propose that \mbox{M~2-9} has descended from a rotating red-giant progenitor and that the 
bipolar morphology is the result of rapid expansion of the ionization front in the N-S directions.

Interferometric mapping of the $^{12}$CO J = 1 - 0 emission from {M~2-9} by Zweigle et al. (1997) indicates 
the existence of an expanding equatorial torus of molecular gas of about 6\arcsec\ diameter around the 
center of {M~2-9}. The symmetry axis of the torus is tilted with respect to the plane of the sky by 
$\sim 17^{\circ}$. The deprojected expansion velocity is 7~km~s$^{-1}$, and for a distance of 650 pc, it 
corresponds to a kinematical age of about 1365 years.

Smith, Balick, \& Gehrz (2005) using high dispersion near infrared spectra observe the 2.12~$\mu$m H$_2$ line and 
determine that the nuclear emission arises in an equatorial molecular disk or torus, consistent with 
the detection of a CO torus by Zweigle et al. (1997), where the H$_2$ emission arises closer to the 
central star, perhaps marking the irradiated inner edge of this larger torus.


%%%%% FIGURA 2

\begin{figure}[!t]
\centering
\includegraphics[width=0.6\columnwidth]{storresfig2.eps}
\caption{Diagram of the structure of {M~2-9} where the knots and nomenclature by Kohoutek 
\& Surdej (1980) are shown. Also indicated is the position of the spectrograph slit used.} 
\label{fm29artist}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% FIGURA 3

\begin{figure}[!t]
\vspace{0.05cm}
\centering
\includegraphics[width=0.99\columnwidth]{storresfig3.eps}
\caption{Position-velocity diagram of the central part of {M~2-9} from the long slit spectrum 
of the H$\alpha$ profile. To determine the systemic velocity it has been assumed that the N3 and 
S3 knots were ejected symmetrically and are moving at the same speed relative to the nucleus. 
The dashed line joins both knots and yields the velocity of the nucleus. The laboratory 
wavelength of the derived velocity of the nucleus is indicated, which corresponds to a heliocentric 
velocity of $+60\pm 5$~km~s$^{-1}$.} 
\label{fm29posvel}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Near-infrared color images of {M~2-9} by Aspin, McLean, \& Smith (1988) suggest that the central core of {M~2-9} is 
surrounded by a dense extended equatorial disk of luminous dust. From the spectral energy distribution 
of the bright central core, using a 4\arcsec\ diameter aperture, Smith \& Gehrz (2005) fit a 260 K 
gray body to the unresolved central engine, although they suggest that the observed SED is actually 
flatter than any single temperature, perhaps indicating dust at various radii in a circumstellar disk.


\subsection{Systemic velocity} 
\label{sm29system}

The determination of the systemic velocity is not straightforward, since no photospheric lines are detected 
in the spectra of the nucleus of {M~2-9}. The emission lines show composite profiles or have 
different velocities depending on the degree of ionization, which allows us to assume that they are produced 
in circumstellar regions. Therefore, we determined the nuclear velocity from the kinematics of the lobes.

We show in Figure~\ref{fm29artist} the positions of the slit in the case of {M~2-9}, where we use the nomenclature 
by Kohoutek \& Surdej (1980) to refer to the condensations. The detailed H$\alpha$ profile was used to determine 
the systemic velocity.

On the long slit H$\alpha$ spectral isocontours shown in Figure~\ref{fm29posvel}, a very bright central region 
with a double peak profile and two regions can be seen that correspond to the northern and southern lobes. 
Superimposed on the lobe emission, there are strong emissions at +15\arcsec\ and --15\arcsec\ relative to the 
nucleus, which correspond to the N3 and S3 knots, respectively. The nuclear velocity was estimated under the 
assumption that both knots were ejected symmetrically from the central region at the same speed. 

In this figure the dashed line joins the intensity maxima of the N3 and S3 knots through the central region. 
The arrows indicate $\lambda_0$, the laboratory wavelength for H$\alpha$ and the central part of the nebula. 
The velocity difference is $+40 \pm 5$~km~s$^{-1}$, which yields a systemic heliocentric velocity of 
$+60 \pm 5$~km~s$^{-1}$ (corresponding to $V_\mathrm{LSR}=+75$~km~s$^{-1}$). In the following sections all 
velocities will be referred to the systemic velocity. This value is comparable to those obtained by other 
authors; namely, $V_\mathrm{hel}=+69.2$~km~s$^{-1}$ from [Fe II] (Smith et al. 2005); and $V_\mathrm{LSR}=+80$~km~s$^{-1}$ 
from CO, (Bachiller, Mart\'in-Pintado, \& Bujarrabal 1990; Zweigle et al. 1997).


%%%%% FIGURA 4

\begin{figure*}[!t]
\centering
\includegraphics[width=14cm]{storresfig4.eps}
\caption{Blue spectrum of the nucleus of {M~2-9}; it has been expanded to show the 
faint lines. The identified lines are indicated.}
\label{fm29blue}
\vspace{0.5cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% FIGURA 5

\begin{figure*}[!t]
\centering
\includegraphics[width=13.8cm]{storresfig5.eps}
\caption{Red spectrum of the nucleus of {M~2-9}; it has been expanded to show the faint 
lines. Also indicated are the identified lines. Our echelle configuration does not cover 
all the wavelength range, and thus there are gaps in the bottom panel.}
\label{fm29red}
\vspace{0.4cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% TABLA 1

\setcounter{table}{0}
\begin{table*}[!t]
\centering
\caption{Emission lines in {M~2-9} -- blue spectrum}
\label{tm29blue}
\setlength{\tabnotewidth}{0.99\textwidth}
\tablecols{13}
\setlength{\tabcolsep}{0.7\tabcolsep}
%\renewcommand{\arraystretch}{0.99999999999}
\casitiny
\begin{tabular}{ccccclccccccl}
\toprule
 &  & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{a}  & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{b}  & FWHM  &  &  &  &  
& I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{a}  & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{b}  & FWHM  & \\ 
$\lambda_0$ & $\lambda_\mathrm{obs}$ &  observed  & ext correct  & \AA  & \multicolumn{1}{c}{Identification}  &  & $\lambda_0$ 
& $\lambda_\mathrm{obs}$ &  observed  & ext correct  & \AA  & \multicolumn{1}{c}{Identification} \\
\cmidrule{1-6}
\cmidrule{8-13} 
3721.94 & 3722.04 & 0.039 & 0.068 & 0.9 & H 14                      &  & 4921.96 & 4922.42 & 0.028 & 0.028 & 0.9 & He I \\
3726.03 & 3726.62 & 0.058 & 0.101 & 0.8 & [O II] 1F                 &  & 4923.93 & 4924.48 & 0.041 & 0.040 & 0.8 & Fe II 42 \\
3728.81 & 3729.09 & 0.040 & 0.069 & 1.0 & [O II] 1F                 &  & 4930.54 & 4931.13 & 0.026 & 0.025 & 1.1 & [Fe III] 1F \\
3734.37 & 3734.76 & 0.013 & 0.023 & 0.9 & H 13                      &  & 4947.37 & 4947.95 & 0.011 & 0.010 & 2.6 & [Fe II] 20F \\
3750.15 & 3750.45 & 0.062 & 0.107 & 1.2 & H 12                      &  & 4950.74 & 4951.41 & 0.019 & 0.018 & 1.1 & [Fe II] 20F \\
3770.63 & 3770.87 & 0.045 & 0.077 & 0.8 & H 11                      &  & 4958.91 & -       & 0.833 & 0.795 & -   & [O III]1F\tabnotemark{\dag}\\
3797.90 & 3798.29 & 0.044 & 0.074 & 0.8 & H 10                      &  & 4973.39 & 4973.05 & 0.007 & 0.006 & 0.3 & [Fe II] 20F \\
3835.38 & 3835.75 & 0.051 & 0.086 & 1.2 & H 9                       &  & 5006.84 & -       & 2.879 & 2.684 & -   & [O III] 1F\tabnotemark{\dag}\\
3863.05 & 3863.22 & 0.013 & 0.022 & 0.8 & Ne II                     &  & 5011.25 & 5011.76 & 0.072 & 0.067 & 1.2 & [Fe III] 1F \\
3869.06 & 3868.96 & 0.232 & 0.384 & 0.9 & [Ne III] 1F               &  & 5015.68 & 5016.13 & 0.057 & 0.053 & 1.1 & He I \\
3889.05 & 3889.25 & 0.183 & 0.300 & 1.2 & H8                        &  & 5019.29 & 5019.12 & 0.017 & 0.016 & 0.7 & O I 15 \\
3964.73 & 3965.29 & 0.012 & 0.020 & 0.4 & He I                      &  & 5020.23 & 5020.80 & 0.025 & 0.023 & 0.8 & [Fe II] 20F \\
3967.79 & 3967.69 & 0.070 & 0.112 & 1.1 & [Ne III] 1F               &  & 5032.70  & 5033.08 & 0.025 & 0.023 & -   & [Fe III] 2F \\
3970.07 & 3970.49 & 0.115 & 0.183 & 1.1 & H$\epsilon$               &  & 5033.59 & 5033.93 & 0.010 & 0.009 & -   & [Fe IV]-F \\
4026.20 & 4026.96 & 0.029 & 0.045 & 0.7 & He I                      &  & 5041.56 & \multirow{2}{*}{{\normalsize \{}5041.94} & \multirow{2}{*}{0.037} & \multirow{2}{*}{0.034} & \multirow{2}{*}{1.5{\normalsize \}}} & [N IV]-F \\
4068.60 & 4069.02 & 0.038 & 0.057 & 0.9 & [S II] 1F                 &  & 5041.40 &         &       &       &     & N I \\
4076.35 & 4076.67 & 0.019 & 0.029 & 1.2 & [S II] 1F                 &  & 5043.52 & 5044.14 & 0.009 & 0.009 & 0.8 & [Fe II] 20F \\
4101.73 & 4102.24 & 0.230 & 0.343 & 1.1 & H$\delta$                 &  & 5047.74 & 5048.43 & 0.015 & 0.014 & 1.3 & He I \\
4143.76 & 4144.43 & 0.009 & 0.013 & 1.1 & He  I                     &  & 5055.80 & 5056.51 & 0.023 & 0.021 & 1.4 & Fe I \\
4177.20 & 4177.82 & 0.016 & 0.022 & 0.7 & [Fe II] 21F               &  & 5060.34 & 5060.73 & 0.006 & 0.005 & 0.8 & [Fe III] \\
4178.85 & \multirow{2}{*}{{\normalsize \{}4179.43} & \multirow{2}{*}{0.018} & \multirow{2}{*}{0.026} & \multirow{2}{*}{0.6{\normalsize \}}} & N II  &  & 5084.77 & 5085.36 & 0.009 & 0.009 & -   & [Fe III] 1F \\ 
4178.96 &         &       &       &     & [Fe II]                   &  & 5107.95 & 5108.59 & 0.010 & 0.009 & 0.9 & [Fe  II] 18F \\
4233.17 & 4233.63 & 0.017 & 0.024 & 0.6 & Fe II 27                  &  & 5111.63 & 5112.31 & 0.031 & 0.027 & 1.0 & [Fe  II] 19F \\
4243.97 & 4244.68 & 0.067 & 0.093 & 1.4 & [Fe II] 21F               &  & 5147.16 & 5147.32 & 0.016 & 0.014 & 1.0 & [Fe I] 19F \\
4276.83 & 4277.33 & 0.045 & 0.062 & 0.7 & [Fe II] 21F               &  & 5158.77 & 5159.28 & 0.145 & 0.126 & 1.2 & [Fe  II] 19F \\
4287.39 & 4287.93 & 0.111 & 0.151 & 0.8 & [Fe II] 7F                &  & 5163.95 & 5164.57 & 0.022 & 0.019 & 1.0 & [Fe  II] \\
4305.89 & 4306.36 & 0.016 & 0.022 & 1.0 & [Fe II] 21F               &  & 5169.03 & 5169.70 & 0.045 & 0.039 & 0.7 & Fe II 42 \\
4319.62 & 4320.17 & 0.032 & 0.042 & 0.8 & [Fe II] 21F               &  & 5181.95 & 5182.63 & 0.016 & 0.014 & 0.8 & [Fe  II] 18F \\
4326.15 & \multirow{2}{*}{{\normalsize \{}4326.94} & \multirow{2}{*}{0.017} & \multirow{2}{*}{0.023} & \multirow{2}{*}{1.0{\normalsize \}}} & [Ni II] 3F  &  & 5191.82 & \multirow{2}{*}{{\normalsize \{}5192.15} & \multirow{2}{*}{0.015} & \multirow{2}{*}{0.013} & \multirow{2}{*}{1.1{\normalsize \}}} & [Ar III] 3F \\
4326.24 &         &       &       &     & [Ni III]                  &  & 5192.01 &         &       &       &     & Ni III \\
4340.47 & 4326.94 & 0.416 & 0.531 &  -  & H$\gamma$\tabnotemark{\dag}  &  & 5197.90 & 5198.34 & 0.018 & 0.015 & 0.6 & [N I] 1F \\
4352.78 & 4352.21 & 0.027 & 0.036 & 1.2 & [Fe  II] 21F              &  & 5199.18 & 5199.77 & 0.011 & 0.010 & 1.0 & [Fe  II] 35F? \\
4359.33 & 4359.69 & 0.099 & 0.129 & 1.5 & [Fe  II] 7F               &  & 5200.26 & 5200.84 & 0.009 & 0.008 & 0.6 & [N I] 1F \\
4363.21 & 4363.55 & 0.361 & 0.468 & 0.8 & [O III] 3F                &  & 5220.06 & 5220.80 & 0.024 & 0.020 & 0.8 & [Fe  II] 19F \\
4372.43 & 4372.92 & 0.011 & 0.014 & 1.2 & [Fe II] 21F               &  & 5233.76 & 5234.24 & 0.014 & 0.012 & 0.8 & [Fe IV]-F \\
4413.78 & 4414.36 & 0.055 & 0.070 & 0.9 & [Fe II] 7F                &  & 5234.62 & 5235.37 & 0.015 & 0.013 & 0.8 & Fe  II 49? \\
4416.27 & 4416.81 & 0.057 & 0.072 & 0.8 & [Fe II] 6F                &  & 5261.62 & 5262.30 & 0.136 & 0.113 & 1.1 & [Fe  II] 19F \\
4452.10 & 4452.61 & 0.037 & 0.046 & 0.8 & [Fe II] 7F                &  & 5268.87 & 5269.50 & 0.028 & 0.023 & 0.5 & [Fe  II] 18F \\
4457.95 & 4458.59 & 0.007 & 0.009 & 0.9 & [Fe II] 6F                &  & 5270.40 & 5270.95 & 0.189 & 0.156 & 1.2 & [Fe III] 1F \\
4471.50 & 4471.86 & 0.062 & 0.075 & 0.8 & He I                      &  & 5273.35 & 5274.00 & 0.106 & 0.088 & 1.0 & [Fe  II] 18F \\
4474.90 & 4475.43 & 0.017 & 0.021 & 0.9 & [Fe II]                   &  & 5275.99 & 5276.66 & 0.041 & 0.034 & 0.8 & Fe  II 49 \\
4485.21 & \multirow{2}{*}{{\normalsize \{}4486.11} & \multirow{2}{*}{0.006} & \multirow{2}{*}{0.008} & \multirow{2}{*}{1.0{\normalsize \}}} & [Ni II] 3F   &  & 5296.83 & 5297.56 & 0.021 & 0.017 & 1.2 & [Fe  II] 19F \\
4485.57 &         &       &       &     & Fe III                    &  & 5333.65 & 5334.38 & 0.069 & 0.055 & 1.0 & [Fe  II] 19F \\
4488.75 & 4489.39 & 0.018 & 0.021 & 1.5 & [Fe II] 6F                &  & 5347.65 & 5348.33 & 0.007 & 0.005 & 0.8 & [Fe  II] 18F \\
4492.63 & \multirow{2}{*}{{\normalsize \{}4492.08} & \multirow{2}{*}{0.008} & \multirow{2}{*}{0.010} & \multirow{2}{*}{0.6{\normalsize \}}} & [Fe II] 6F   &  & 5363.30  & 5363.78 & 0.014 & 0.011 & 1.1 & [Ni IV]-F \\
4492.41 &         &       &       &     & Ne I                      &  & 5376.45 & 5377.21 & 0.092 & 0.073 & 0.9 & [Fe  II] 19F \\
4509.61 & 4509.47 & 0.008 & 0.010 & 1.6 & [Fe II] 6F                &  & 5412.64 & 5413.08 & 0.044 & 0.034 & 1.4 & [Fe  II] 17F \\
4515.34 & 4515.70 & 0.020 & 0.024 & 1.0 & Fe II                     &  & 5433.13 & 5433.83 & 0.026 & 0.020 & 0.9 & [Fe  II] 18F \\
4583.84 & 4584.42 & 0.029 & 0.033 & 0.5 & Fe II 38                  &  & 5477.24 & 5478.10 & 0.013 & 0.010 & 1.2 & [Fe  II] 34F \\
4607.03 & 4607.57 & 0.022 & 0.025 & 1.5 & [Fe III] 3F               &  & 5495.82 & 5496.56 & 0.016 & 0.012 & 1.7 & [Fe  II] 17F \\
4629.39 & 4629.95 & 0.016 & 0.018 & 0.9 & [Fe II] 37                &  & 5527.33 & 5528.04 & 0.065 & 0.049 & 0.9 & [Fe  II] 17F \\
4639.67 & 4640.29 & 0.016 & 0.018 & 0.9 & [Fe II] 4                 &  & 5534.89 & \multirow{2}{*}{{\normalsize \{}5535.55} & \multirow{2}{*}{0.014} & \multirow{2}{*}{0.011} & \multirow{2}{*}{0.9{\normalsize \}}} & Fe  II  \\
4658.05 & 4658.56 & 0.136 & 0.150 & 1.0 & [Fe III] 3F               &  & 5534.70 &         &       &       &     & Ni II \\
4701.53 & 4702.02 & 0.117 & 0.126 & 0.5 & [Fe III] 3F               &  & 5551.91 & \multirow{2}{*}{{\normalsize \{}5552.38} & \multirow{2}{*}{0.008} & \multirow{2}{*}{0.006} & \multirow{2}{*}{1.7{\normalsize \}}} & [Mn  II] \\
4713.14 & 4713.61 & 0.017 & 0.018 & 1.0 & He I                      &  & 5551.96 &         &       &       &     & [Cr III] \\
4728.07 & 4728.61 & 0.037 & 0.040 & 0.9 & [Fe II] 4F                &  & 5556.29 & 5556.54 & 0.008 & 0.006 & 2.3 & [Fe  II] 18F \\
4733.91 & 4734.28 & 0.024 & 0.026 & 1.3 & [Fe III] 3F               &  & 5580.82 & 5581.52 & 0.004 & 0.003 & 1.3 & [Fe  II] 39F \\
4754.69 & 4755.27 & 0.024 & 0.025 & 1.0 & [Fe III] 3F               &  & 5613.27 & 5614.25 & 0.008 & 0.006 & 2.0 & [Fe  II] 39F \\
4769.43 & 4769.91 & 0.020 & 0.021 & 1.2 & [Fe III] 3F               &  & 5673.21 & 5674.13 & 0.014 & 0.010 & 1.1 & [Fe II] \\
4774.72 & 4775.34 & 0.012 & 0.012 & 1.1 & [Fe II] 20F               &  & 5746.96 & 5747.75 & 0.035 & 0.024 & 1.4 & [Fe  II] 34F \\
4777.68 & 4778.19 & 0.012 & 0.012 & 1.0 & [Fe III] 3F               &  & 5754.64 & 5755.24 & 0.105 & 0.072 & 1.2 & [N  II] 3F \\
4814.53 & 4815.11 & 0.080 & 0.082 & 0.9 & [Fe II] 20F               &  & 5773.40 & 5774.06 & 0.004 & 0.002 & 0.7 & [Fe I] 2F? \\
4861.33 & -       & 1.000 & 1.000 & -   & H$\beta$\tabnotemark{\dag}   &  & 5780.20 & \multirow{2}{*}{{\normalsize \{}5780.72} & \multirow{2}{*}{0.008} & \multirow{2}{*}{0.005} & \multirow{2}{*}{1.7{\normalsize \}}} & Fe II \\
4867.95 & \multirow{2}{*}{{\normalsize \{}4868.12} & \multirow{2}{*}{0.008} & \multirow{2}{*}{0.008} & \multirow{2}{*}{1.0{\normalsize \}}} & [Fe IV]-F   &  & 5780.51 &         &       &       &     & Ne II \\ 
4868.00 &         &       &       &     & Ne II                     &  & 5800.70  & 5800.81 & 0.012 & 0.008 & 1.6 & [Fe IV]-F \\ 
4874.49 & 4875.12 & 0.022 & 0.021 & 1.0 & [Fe II] 20F               &  & 5816.02 & \multirow{3}{*}{\big\{5816.75} & \multirow{3}{*}{0.007} & \multirow{3}{*}{0.005} & \multirow{3}{*}{1.0\big\}} & S II \\
4881.00 & 4881.76 & 0.018 & 0.018 & 2.2 & [Fe III] 2F               &  & 5816.00 &         &       &       &     & Mg I \\
4900.50 & 4899.21 & 0.007 & 0.007 & 0.8 & [Fe IV]-F                 &  & 5816.08 &         &       &       &     & Si I \\ 
4905.81 & 4905.86 & 0.070 & 0.068 & 0.2 & [Fe II] 20F               &  & 5835.45 & 5836.13 & 0.012 & 0.008 & 1.0 & [Fe  II] \\ 
4906.36 & \multirow{3}{*}{\big\{4906.94} & \multirow{3}{*}{0.029} & \multirow{3}{*}{0.029} & \multirow{3}{*}{0.7\big\}} & O II 28   &  & 5870.02 & 5870.81 & 0.005 & 0.003 & 1.4 & [Fe  II] \\
4905.34 &         &       &       &     & [Fe III]                  &  & 5875.67 & 5876.29 & 0.595 & 0.391 & 1.3 & He I 11 \\
4906.56 &         &       &       &     & [Fe IV]                   &  & 5914.61 & 5914.82 & 0.004 & 0.003 & - & Ne II \\
\bottomrule
%\cmidrule{1-6}
%\cmidrule{8-13}
\end{tabular}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%% TABLA 1 continua

\setcounter{table}{0}
\begin{table*}[!t]
\vspace{0.1cm}
\centering
\centerline{TABLE~1 (CONTINUED)\smallskip}
\setlength{\tabnotewidth}{0.99\textwidth}
\tablecols{13}
\setlength{\tabcolsep}{0.68\tabcolsep}
\casitiny
\begin{tabular}{ccccclccccccl}
\toprule
  &  & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{a}  & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{b}  & FWHM   &  &  &  &  
& I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{a}  & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{b}  & FWHM  & \\
$\lambda_0$  & $\lambda_\mathrm{obs}$  & observed  & ext correct  &\AA  & Identification  &  & $\lambda_0$  & $\lambda_\mathrm{obs}$  
& observed  & ext correct  &\AA  & Identification \\ 
\cmidrule{1-6}
\cmidrule{8-13} 
5957.56 & 5958.54 & 0.063 & 0.040 & 2.2~ & Si  II  4    &  & 6383.72 & 6384.66 & ~0.009 & ~0.005 & 0.8~ & Fe  II \\
5978.93 & 5979.64 & 0.059 & 0.037 & 1.4~ & Si  II 4     &  & 6385.45 & \multirow{2}{*}{{\normalsize \{}6386.25} & \multirow{2}{*}{~0.008} & \multirow{2}{*}{~0.004} & \multirow{2}{*}{0.6{\normalsize \}}} & Fe  II \\
6000.2~ & 6000.60 & 0.008 & 0.005 & 2.0~ & [Ni III] 2F  &  & 6385.64 &         &       &       &     & O I \\
6046.44 & 6047.14 & 0.039 & 0.024 & 1.3~ & O I  22      &  & 6440.40  & 6441.02 & ~0.007 & ~0.004 & 1.2~ & [Fe II] 15F \\
6256.71 & \multirow{3}{*}{\big\{6257.68} & \multirow{3}{*}{0.007} & \multirow{3}{*}{0.004} & \multirow{3}{*}{1.3\big\}} & Si I   &  & 6442.95 & 6443.82 & ~0.005 & ~0.003 & 0.98 & Fe II \\
6256.93 &         &       &       &     & O I          &  & 6456.38 & 6457.31 & ~0.018 & ~0.010 & 1.1~ & [Fe II] \\
6257.20 &         &       &       &     & C II         &  & 6491.28 & 6492.33 & ~0.023 & ~0.012 & -   & Fe II \\
6300.30 & 6301.10 & 0.620 & 0.350 & 1.0~ & [O I] 1F     &  & 6548.04 & 6548.82 & ~0.187 & ~0.097 & 2.0~ & [N II] 1F \\
6312.06 & 6312.68 & 0.374 & 0.211 & 1.2~ & [S III]      &  & 6562.80 & -       & 33.636 & 17.380 & -   & H$\alpha$\tabnotemark{\dag} \\
6317.98 & 6318.88 & 0.032 & 0.018 & 1.1~ & Fe  II       &  & 6583.46 & 6584.20 & ~0.541 & ~0.277 & 2.0~ & [N II] \\
6347.11 & 6347.87 & 0.066 & 0.037 & 1.6~ & Si  II (2)   &  & 6666.80 & 6667.60 & ~0.034 & ~0.017 & 0.7~ & [Ni II] (2D-2F) \\
6363.78 & 6364.57 & 0.166 & 0.092 & 0.9~ & [O I] 1F     &  & 6678.15 & 6678.77 & ~0.303 & ~0.151 & 1.4~ & He I (46) \\
6371.37 & \multirow{2}{*}{{\normalsize \{}6372.09} & \multirow{2}{*}{0.028} & \multirow{2}{*}{0.015} & \multirow{2}{*}{1.2{\normalsize \}}} & Si  II (2)   &  & 6700.64 & \multirow{2}{*}{{\normalsize \{}6702.94} & \multirow{2}{*}{~0.010} & \multirow{2}{*}{~0.005} & \multirow{2}{*}{0.6{\normalsize \}}} & [Fe II] \\
6371.13 &         &       &       &     & Fe II        &  & 6702.22 &         &       &       &     & C II \\
6379.46 & \multirow{2}{*}{{\normalsize \{}6380.14} & \multirow{2}{*}{0.010} & \multirow{2}{*}{0.005} & \multirow{2}{*}{0.73{\normalsize \}}} & Ni  II  &  & 6716.44 & 6717.47 & ~0.021 & ~0.010 & 1.7~ & [S II] 2F \\
6380.16 &         &       &       &      & O II        &  & 6730.82 & 6731.75 & ~0.046 & ~0.022 & 1.7~ & [S II] 2F \\
\bottomrule
\noalign{\smallskip}
\tabnotetext{a}{I(H$\beta$)$_\mathrm{obs}=1.50 \times 10^{-13}$~erg~cm$^{-2}$~s$^{-1}$.}
\tabnotetext{b}{I(H$\beta$)$_\mathrm{extcorr}=2.35 \times 10^{-12}$~erg~cm$^{-2}$~s$^{-1}$.} 
\tabnotetext{\dag}{Several components.} 
\end{tabular}
\vspace{0.3cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% TABLA 2

\setcounter{table}{1}
\begin{table*}[!t]
\vspace{0.1cm}
\centering
\caption{Emission lines in {M~2-9} -- red spectrum}
\label{tm29red}
\setlength{\tabnotewidth}{0.99\textwidth}
\tablecols{13}
\setlength{\tabcolsep}{0.33\tabcolsep}
\tiny
\begin{tabular}{ccccclccccccl}
\toprule
  &   & $100 \times$ I$_\lambda$/I$_\mathrm{H\alpha}$\tabnotemark{a}  & $100 \times$ I$_\lambda$/I$_\mathrm{H\alpha}$\tabnotemark{b}  & FWHM  
&   &   &   &   & $100 \times$ I$_\lambda$/I$_\mathrm{H\alpha}$\tabnotemark{a}  & $100 \times$ I$_\lambda$/I$_\mathrm{H\alpha}$\tabnotemark{b}  
& FWHM  &  \\
$\lambda_0$  & $\lambda_\mathrm{obs}$  & observed  & ext corrected  & \AA  & \multicolumn{1}{c}{Identification}  &  &  $\lambda_0$  & $\lambda_\mathrm{obs}$  & observed  & ext corrected  & \AA  & \multicolumn{1}{c}{Identification} \\  
\cmidrule{1-6}
\cmidrule{8-13} 
6548.05 & 6548.99 & ~~0.571 & ~~0.573 & 1.8~ & [N II] 1F                 &  & 8323.42 & 8324.26 & 0.076 & 0.044 & 1.5~ & H I Pa25 \\
6562.82 & -       & 100.000 & 100.000 & -    & H$\alpha$\tabnotemark{*}  &  & 8330.59 & 8330.92 & 0.029 & 0.017 & 1.4~ & Fe II \\
6583.46 & 6584.24 & ~~1.747 & ~~1.735 & 1.9~ & [N II] 1F                 &  & 8333.78 & 8334.50 & 0.029 & 0.017 & 1.2~ & H I Pa24 \\ 
6666.80 & 6668.05 & ~~0.104 & ~~0.100 & 1.5~ & [Ni II] (2D-2F)           &  & 8357.69 & 8359.04 & 0.175 & 0.099 & 2.7~ & [Cr  II](1F) + H I Pa22 \\
6678.15 & 6678.99 & ~~0.765 & ~~0.736 & 1.4~ & He I (46)                 &  & 8413.32 & 8414.73 & 0.124 & 0.069 & 1.5~ & H I Pa19 \\
6716.44 & 6717.65 & ~~0.068 & ~~0.064 & 1.6~ & [S II] 2F                 &  & 8437.95 & 8439.88 & 0.114 & 0.063 & 1.85 & H I Pa18 \\
6730.81 & 6731.81 & ~~0.204 & ~~0.193 & 2.1~ & [S II] 2F                 &  & 8446.48 & \multirow{2}{*}{{\normalsize \{}8448.25} & \multirow{2}{*}{4.392} & \multirow{2}{*}{2.430} & \multirow{2}{*}{1.6{\normalsize \}}} & O I (4)? \\
6746.73 & \multirow{2}{*}{{\normalsize \{}6747.81} & \multirow{2}{*}{~~0.020} & \multirow{2}{*}{~~0.018} & \multirow{2}{*}{1.0{\normalsize \}}} & [Fe II] (4P-2S)           &  & 8446.12 &         &       &       &      & Fe II \\
6746.13 &         &         &         &      & [Fe II]                   &  & 8467.26 & 8469.90 & 0.510 & 0.280 & 2.7~ & H I Pa17 \\
6997.10 & \multirow{2}{*}{{\normalsize \{}6997.88} & \multirow{2}{*}{~~0.094} & \multirow{2}{*}{~~0.082} & \multirow{2}{*}{1.2{\normalsize \}}} & [Fe IV] 4P                &  & 8490.73 & 8491.73 & 0.525 & 0.287 & 1.5~ & [Fe  II] (2G-2F) \\
6997.18 &         &         &         &      & Fe I]                     &  & 8499.70 & \multirow{3}{*}{\big\{8501.24} & \multirow{3}{*}{0.376} & \multirow{3}{*}{0.205} & \multirow{3}{*}{1.7\big\}} & [Ni III] \\ 
7002.12 & 7003.07 & ~~0.149 & ~~0.129 & 1.6~ & O I (21)                  &  & 8500.20 &         &       &       &      & [Cl III] \\
7065.18 & 7066.10 & ~~0.933 & ~~0.792 & 1.4~ & He I (10)                 &  & 8500.32 &         &       &       &      & Cl III \\
7131.11 & 7131.71 & ~~0.041 & ~~0.034 & 1.2~ & [Fe II]                   &  & 8509.66 & 8509.92 & 0.035 & 0.019 & 1.23 & Fe II \\
7135.80 & 7136.70 & ~~1.816 & ~~1.506 & 1.8~ & [Ar III] (1F)             &  & 8523.66 & 8525.14 & 0.047 & 0.025 & 1.37 & Fe II] \\ 
7155.16 & 7156.37 & ~~0.578 & ~~0.477 & 1.3~ & [Fe II] (14F)             &  & 8545.38 & 8546.67 & 0.490 & 0.264 & 2.0~ & H I Pa15 \\
7169.24 & 7169.78 & ~~0.051 & ~~0.042 & 2.1~ & Fe II                     &  & 8578.69 & \multirow{2}{*}{{\normalsize \{}8579.84} & \multirow{2}{*}{0.075} & \multirow{2}{*}{0.040} & \multirow{2}{*}{1.6{\normalsize \}}} & [Cl  II] (1F) \\
7172.00 & 7173.13 & ~~0.233 & ~~0.191 & 1.6~ & [Fe II] (14F)             &  & 8579.18 &         &       &       &      & Fe I \\  
7182.09 & 7184.26 & ~~0.053 & ~~0.043 & 2.5~ & O II                      &  & 8598.39 & 8599.48 & 0.431 & 0.229 & 1.8~ & H I Pa14 \\
7203.90 & \multirow{2}{*}{{\normalsize \{}7203.96} & \multirow{2}{*}{~~0.033} & \multirow{2}{*}{~~0.027} & \multirow{2}{*}{1.8{\normalsize \}}} & N II                      &  & 8616.95 & 8618.22 & 0.510 & 0.269 & 1.5~ & [Fe  II] (13F) \\
7203.35 &         &         &         &      & Fe I                      &  & 8626.40 & 8626.05 & 0.027 & 0.014 & 0.7~ & [Cr II] \\
7222.18 & \multirow{2}{*}{{\normalsize \{}7223.69} & \multirow{2}{*}{~~0.029} & \multirow{2}{*}{~~0.024} & \multirow{2}{*}{2.1{\normalsize \}}} & [Fe II]                   &  & 8665.02 & 8666.07 & 0.343 & 0.179 & 1.9~ & H I Pa13 \\
7222.80 &         &         &         &      & [Fe IV] (4D-4F)           &  & 8672.32 & \multirow{2}{*}{{\normalsize \{}8673.60} & \multirow{2}{*}{0.184} & \multirow{2}{*}{0.096} & \multirow{2}{*}{1.73{\normalsize \}}} & Fe II] \\ 
7236.19 & \multirow{2}{*}{{\normalsize \{}7237.71} & \multirow{2}{*}{~~0.061} & \multirow{2}{*}{~~0.049} & \multirow{2}{*}{1.8{\normalsize \}}} & C I                       &  & 8672.40 &         &       &       &      & O III \\
7236.52 &         &         &         &      & Fe I                      &  & 8680.20 & 8681.34 & 0.067 & 0.035 & 2.0~ & O II \\
7254.44 & 7255.43 & ~~0.288 & ~~0.230 & 1.9~ & O I(20)                   &  & 8703.87 & 8704.61 & 0.082 & 0.042 & 1.9~ & [Ni II] \\
7281.35 & \multirow{2}{*}{{\normalsize \{}7281.83} & \multirow{2}{*}{~~0.178} & \multirow{2}{*}{~~0.141} & \multirow{2}{*}{1.6{\normalsize \}}} & He I (45)                 &  & 8722.35 & 8723.35 & 0.067 & 0.034 & 1.02 & Fe II \\
7281.63 &         &         &         &      & [Fe II]                   &  & 8806.76 & \multirow{3}{*}{\big\{8807.32} & \multirow{3}{*}{0.131} & \multirow{3}{*}{0.066} & \multirow{3}{*}{1.6\big\}} & Mg I (7) \\ 
7131.12 & 7131.71 & ~~0.041 & ~~0.034 & 1.2~ & [Fe II] (2G-2G)           &  & 8805.43 &         &       &       &      & Fe I \\
7319.99 & 7321.14 & ~~2.471 & ~~1.931 & 2.0~ & [O II] (2F)               &  & 8806.06 &         &       &       &      & Fe I \\
7330.73 & 7331.44 & ~~2.510 & ~~1.955 & 2.3~ & [O II] (2F)               &  & 8830.50 & \multirow{2}{*}{{\normalsize \{}8831.55} & \multirow{2}{*}{0.157} & \multirow{2}{*}{0.078} & \multirow{2}{*}{1.9{\normalsize \}}} & He I \\
7377.83 & 7379.16 & ~~0.494 & ~~0.379 & 1.8~ & [Ni II] (2F)              &  & 8830.70 &         &       &       &      & [Fe III] \\
7388.17 & 7389.47 & ~~0.175 & ~~0.133 & 1.5~ & [Fe II] (14F)             &  & 8862.79 & 8864.18 & 0.827 & 0.409 & 1.9~ & H I Pa11 \\
7411.61 & 7412.79 & ~~0.208 & ~~0.158 & 1.6~ & [Ni II] (2F)              &  & 8876.34 & 8877.52 & 0.086 & 0.043 & 1.8~ & Ca I  \\
7423.64 & 7424.97 & ~~0.043 & ~~0.033 & 1.6~ & N I (3)                   &  & 8885.62 & \multirow{4}{*}{\Big\{8893.48} & \multirow{4}{*}{0.194} & \multirow{4}{*}{0.095} & \multirow{4}{*}{1.8\Big\}} & [Fe  II]  \\
7442.30 & 7443.42 & ~~0.047 & ~~0.035 & 1.2~ & N I (3)                   &  & 8884.94 &         &       &       &      & Si I \\ 
7452.54 & 7453.68 & ~~0.188 & ~~0.141 & 1.2~ & [Fe II] (14F)             &  & 8890.90 &         &       &       &      & Fe  II  \\
7468.31 & 7469.36 & ~~0.061 & ~~0.045 & 1.2~ & N I (3)                   &  & 8891.91 &         &       &       &      & [Fe  II]  \\
7513.16 & 7514.25 & ~~0.043 & ~~0.032 & 1.72 & Fe II                     &  & 8912.63 & \multirow{2}{*}{{\normalsize \{}8913.84} & \multirow{2}{*}{0.076} & \multirow{2}{*}{0.037} & \multirow{2}{*}{1.2{\normalsize \}}} & Fe II] \\
7513.21 & 7514.25 & ~~0.175 & ~~0.133 & 1.5~ & Ca I                      &  & 8913.27 &         &       &       &      & Mg II \\    
7637.54 & 7638.63 & ~~0.067 & ~~0.047 & 1.5~ & [Fe  II] (1F)             &  & 8916.76 & 8917.78 & 0.147 & 0.072 & 1.0~ & Fe I \\
7686.94 & 7688.06 & ~~0.067 & ~~0.046 & 1.7~ & [Fe  II] (1F)             &  & 8926.64 & \multirow{2}{*}{{\normalsize \{}8927.97} & \multirow{2}{*}{1.229} & \multirow{2}{*}{0.598} & \multirow{2}{*}{1.51{\normalsize \}}} & Fe II \\
7711.43 & 7712.52 & ~~0.202 & ~~0.139 & 1.1~ & Fe II                     &  & 8927.40 &         &       &       &      & Ne I \\ 
7751.12 & 7751.58 & ~~0.484 & ~~0.330 & 2.4~ & [Ar III] (1F)             &  & 9014.91 & 9015.80 & 0.471 & 0.224 & 1.7~ & H I Pa10 \\
7866.55 & 7867.97 & ~~0.204 & ~~0.134 & 1.60 & Fe II                     &  & 9033.50 & \multirow{2}{*}{{\normalsize \{}9034.82} & \multirow{2}{*}{0.076} & \multirow{2}{*}{0.036} & \multirow{2}{*}{1.5{\normalsize \}}} & [Fe  II] (13F) \\
7874.25 & \multirow{2}{*}{{\normalsize \{}7878.68} & \multirow{2}{*}{~~0.161} & \multirow{2}{*}{~~0.105} & \multirow{2}{*}{2.8{\normalsize \}}} & [Fe II]                   &  & 9034.20 &         &       &       &      & C II \\ 
7878.05 &         &         &         &      &  Mg II                    &  & 9051.95 & 9053.66 & 0.118 & 0.055 & 1.7~ & [Fe  II] (13F) \\
7889.90 & 7890.95 & ~~0.255 & ~~0.166 & 2.7~ & [Ni III] (1F)             &  & 9068.9~ & 9070.08 & 1.961 & 0.921 & 2.3~ & [S III] (1F) \\
7896.29 & 7897.31 & ~~0.167 & ~~0.109 & 2.9~ & [Fe II]                   &  & 9177.13 & \multirow{2}{*}{{\normalsize \{}9177.57} & \multirow{2}{*}{0.765} & \multirow{2}{*}{0.350} & \multirow{2}{*}{1.7{\normalsize \}}} & Fe I \\
7917.80 & 7919.00 & ~~0.110 & ~~0.071 & 1.9~ & Fe II                     &  & 9178.42 &         &       &       &      & C I \\ 
7971.6~ & \multirow{2}{*}{{\normalsize \{}7971.51} & \multirow{2}{*}{~~0.071} & \multirow{2}{*}{~~0.045} & \multirow{2}{*}{1.4{\normalsize \}}} & He I                      &  & 9178.05 & 9179.52 & 0.784 & 0.359 & 2.4~ & Fe II \\
7970.61 &         &         &         &      & C III                     &  & 9188.7~ & \multirow{2}{*}{{\normalsize \{}9189.64} & \multirow{2}{*}{0.127} & \multirow{2}{*}{0.058} & \multirow{2}{*}{1.6{\normalsize \}}} & Ne II \\
8000.08 & \multirow{2}{*}{{\normalsize \{}8001.15} & \multirow{2}{*}{~~0.192} & \multirow{2}{*}{~~0.121} & \multirow{2}{*}{1.5{\normalsize \}}} & [Cr  II] (1F)             &  & 9188.16 &         &       &       &      & Fe II] \\ 
7999.49 &         &         &         &      & [Fe II]                   &  & 9197.9~ & 9198.72 & 0.431 & 0.196 & 1.9~ & C II \\ 
8125.30 & 8126.87 & ~~0.061 & ~~0.037 & 1.4~ & [Cr  II] (1F)             &  & 9205.30 & 9205.08 & 0.665 & 0.302 & 1.4~ & C II \\
8158.13 & 8159.24 & ~~0.098 & ~~0.059 & 1.70 & Fe I                      &  & 9210.24 & \multirow{2}{*}{{\normalsize \{}9212.54} & \multirow{2}{*}{0.188} & \multirow{2}{*}{0.085} & \multirow{2}{*}{2.8{\normalsize \}}} & He I (6/9) \\
8223.13 & \multirow{2}{*}{{\normalsize \{}8224.49} & \multirow{2}{*}{~~0.104} & \multirow{2}{*}{~~0.061} & \multirow{2}{*}{1.6{\normalsize \}}} & N I (2)                   &  & 9210.89 &         &       &       &      & Fe II \\
8222.04 &         &         &         &      & [Fe II]                   &  & 9219.80 & \multirow{3}{*}{\big\{9219.91} & \multirow{3}{*}{0.320} & \multirow{3}{*}{0.145} & \multirow{3}{*}{2.2\big\}} & C II \\  
8229.67 & 8230.35 & ~~0.249 & ~~0.147 & 1.1~ & [Cr  II] (1F)             &  & 9218.92 &        &       &       &      & Fe II \\
8235.35 & \multirow{3}{*}{\big\{8235.68} & \multirow{3}{*}{~~0.098} & \multirow{3}{*}{~~0.058} & \multirow{3}{*}{2.1\big\}} & O I                       &  & 9218.25 &        &       &       &      & Mg II \\
8235.00 &         &         &         &      & Mg I                      &  & 9229.02 & 9230.57 & 1.584 & 0.716 & 2.0~ & Pa9  \\ 
8235.69 &         &         &         &      & Mg II                     &  & 9244.26 & \multirow{2}{*}{{\normalsize \{}9245.77} & \multirow{2}{*}{0.225} & \multirow{2}{*}{0.102} & \multirow{2}{*}{2.3{\normalsize \}}} & Mg II (1) \\
8242.39 & \multirow{2}{*}{{\normalsize \{}8243.58} & \multirow{2}{*}{~~0.047} & \multirow{2}{*}{~~0.028} & \multirow{2}{*}{1.3{\normalsize \}}} & N I (2)                   &  & 9244.74 &         &       &       &      & Fe II \\
8242.44 &         &         &         &      & Fe I                      &  & 9267.56 & 9269.20 & 0.112 & 0.050 & 1.6~ & [Fe  II] (13F) \\ 
8286.43 & 8288.67 & ~~0.273 & ~~0.158 & 1.3~ & H I                       &  & 9297.1~ & \multirow{2}{*}{{\normalsize \{}9298.47} & \multirow{2}{*}{0.431} & \multirow{2}{*}{0.192} & \multirow{2}{*}{1.57{\normalsize \}}} & Fe I \\
8306.11 & 8306.99 & ~~0.098 & ~~0.056 & 1.4~ & H I Pa27                  &  & 9297.1~ &         &       &       &      & Fe II \\
8308.49 & 8309.51 & ~~0.131 & ~~0.076 & 1.1~ & [Cr  II] (1F)             &  & 9388.87 & 9389.99 & 0.333 & 0.145 & 2.8~ & Mn II \\
8314.26 & 8315.35 & ~~0.084 & ~~0.048 & 1.4~ & H I Pa26                  &  & 9407.08 & 9408.22 & 0.378 & 0.164 & 2.1~ & Fe II \\
\bottomrule
\noalign{\smallskip}
\tabnotetext{a}{I(H$\alpha$)$_\mathrm{obs}=6.10 \times 10^{-12}$~erg~cm$^{-2}$~s$^{-1}$.} 
%\textsuperscript{b}Where I(H$\alpha$)$_\mathrm{extcorr}=4.40 \times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$.}
\tabnotetext{b}{I(H$\alpha$)$_\mathrm{extcorr}=4.40 \times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$.}
\tabnotetext{*}{Several components.} 
\end{tabular}
\vspace{0.4cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\subsection{Spectrum Description} 
\label{sm29spectrum}

The nuclear spectrum is dominated by emission lines over a weak continuum, where no absorption lines can be 
detected; therefore it is not possible to determine directly the central star spectral type. In our data, 
emission lines from the following ions were identified and measured: H~I, He~I, C~I, C~II, C~III, N~I, N~II, 
O~I, O~II, O~III, Ne~I, Ne~II, Mg~I, Mg~II, Si~I, Si~II, Cl~III, Ca~I, Fe~I, Fe~II, Fe~III, Ni~II, Ni~III, 
[N~I], [N~II], [N~IV], [O~I], [O~II], [O~III], [Ne~III], [S~II], [S~III], [Cl~II], [Cl~III], [Ar~III], [Cr~II], 
[Cr~III], [Mn~II], [Fe~I], [Fe~II], [Fe~III], [Fe~IV], [Ni~II], [Ni~III], and [Ni~IV]. To identify the lines 
in this very rich spectrum we used the code EMILI (Sharpee et al. 2003) under a set of different temperatures 
(9,000 to 11,500 K) and densities ($10^4$ to $10^7$~cm$^{-3}$) to account for the different conditions that 
prevail in the source. The blue spectrum from 4150 to 6850~\AA\ is presented in Figure~\ref{fm29blue}, and the 
red spectrum from 6500 to 9400~\AA\ in Figure~\ref{fm29red}. In both figures the most intense features have 
been marked. 

Some of these lines are emitted in different unresolved regions of different velocities. For example, the Balmer 
lines have two maxima with a stronger red component. Also the profiles of the [OIII] 4959,5007~\AA\ lines show 
two maxima, but in this case the blue component is brighter. The rest of the lines are approximately symmetric 
about a single peak; at the spectral resolution available, they are adjusted satisfactorily with Gaussian profiles. 
Tables~1 and 2 list the identifications of the observed lines for the 3700 to 6800~\AA\ range and for the 6450 
to 9460~\AA\ range, respectively. They contain in Columns~(1) and (2) the laboratory and observed wavelengths of 
the identified lines; in Columns~(3) and (4) the measured intensities relative to H$\beta$ and the extinction 
corrected intensity (see details of the extinction law in \S~\ref{sm29extinc}); in Column~(5) the measured 
FWHM (uncorrected for instrumental width); and in Column~(6) the line identification. The blue section is normalized 
to H$\beta$, while the red section has been normalized to H$\alpha$. It was not possible to normalize the red 
spectrum to H$\beta$ because the observations were carried out on different nights under different atmospheric 
conditions.


\subsection{Interstellar extinction} 
\label{sm29extinc}

Usually the Balmer and Paschen line ratios are used to derive interstellar extinction to any object. {M~2-9} poses 
a difficult case, since its hydrogen lines show very complex profiles. While the Balmer lines show double peaked 
profiles, the Paschen lines show single component lines. Figure~\ref{fm29hydrogen} shows these profiles normalized 
to maximum intensity. In this figure, lines are shown in increasing order of the upper level of the transition 
for each series. In the Balmer series, the H$\alpha$ line shows two well separated components and the separation 
becomes less well defined as the energy of the upper level increases.

On the other hand, the Paschen series shows essentially a single component. A possible explanation for the single 
components could be that at these wavelengths the spectral resolution is lower than in the case of H$\alpha$. To 
rule out this possibility we degraded the resolution of the Balmer series to the same resolution as the Paschen 
series and found that the difference between the two peaks continues to be appreciable. Alternatively, if extinction 
were responsible for the shape of the Balmer profiles, the effect would tend to depend on the wavelength of the lines, 
and this effect is not observed.

In addition to the observed difference in the shape of the profile of the hydrogen lines, it should be noted that, 
within the uncertainties, in each series, for increasing energy of the upper level the centroid of the line approaches 
the systemic velocity. A possible explanation for this behavior in terms of the expected line ratios for each one 
of the transitions and interstellar extinction will be described in \S~\ref{sm29forbid}.


%%%%% FIGURA 6

\begin{figure}[!t]
\centering
\includegraphics[width=0.98\columnwidth]{storresfig6.eps}
\caption{Balmer and Paschen hydrogen line series in the nucleus of {M~2-9}. The 
intensities have been normalized. H$\alpha$ shows a marked velocity shift, and the 
Balmer line profiles show two components.}
\label{fm29hydrogen}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


We compared the hydrogen emission lines in the nucleus with the predictions for Case B recombination (Osterbrock 1989) 
in an attempt to derive $A_\mathrm{V}$ by means of a normal interstellar extinction law. 

Assuming a mean density and temperature of $10^6$~cm$^{-3}$ and $10^4$~K and adopting the interstellar extinction curve 
given by Cardelli, Clayton, \& Mathis (1989) for $R_\mathrm{V}=3.1$, it is found that while the observed H$\alpha$ to 
H$\beta$ ratio suggests $A_\mathrm{V}=6.45$, all other Balmer lines are consistent with $A_\mathrm{V}=1.3\pm 0.2$. Both 
estimates are quite different from the value determined by Hora \& Latter (1994), $A_\mathrm{V}=2.3$, from the Brackett 
lines in the nearby IR spectrum of the nucleus. The anomalous value of extinction obtained using the H$\alpha$/H$\beta$ 
ratio can be explained as an exceptionally intense H$\alpha$ line.

The explanation of the discrepancy between the values of $A_\mathrm{V}$ determined by the other Balmer lines and the Brackett 
lines is less obvious. This discrepancy exceeds all the reasonable uncertainties in line intensity determinations and the 
possible deviation from Case B recombination. Therefore, this discrepancy suggests an unusual extinction curve.

The grain distribution of a nebular gas under the presence of a strong radiation field is expected to correspond to 
grains of larger sizes than those in the typical interstellar medium. This is due to the fact that the small dust 
grains are easily photodissociated while the large ones are much less affected. A larger grain size results in a 
shallower extinction curve than that of the galactic interstellar medium. A well known example of this effect is the 
extinction curve of the Orion Nebula (Mathis \& Wallenhorst 1981; Kim, Martin, \& Hendry 1994), which is best fitted 
by $R_\mathrm{V}=5.1$. Besides the $R_\mathrm{V}$ variations, the nebula can exhibit an uncommon extinction curve due 
to the reflection effect and to the mixture of dust grains along the emitting gas column. The dependence on wavelength 
of the modified extinction curve can be characterized by the nebular extinction parameter, $N_\mathrm{V}$ (Bautista, 
Pogge, \& DePoy 1995).

An important characteristic of interstellar extinction is that for $\lambda$ longer than $\sim 1~\mu$m, it is practically 
insensitive to $R_\mathrm{V}$ and $N_\mathrm{V}$. Then, the observations of Hora \& Latter (1994) could be real indicators 
of the amount of extinction in the $J$ band. Starting from their value of $A_\mathrm{V}=2.3$ and $E_{B-V}=0.75$, a value of 
$A_\mathrm{J}=0.12$ is obtained. With this fixed value we derive $R_\mathrm{V}=5.0 \pm 0.5$ and $A_\mathrm{V}=2.7 \pm 0.1$ to 
best represent the optical spectrum. These values were used to correct the observed nuclear lines in our spectrum (Tables~1 
and 2). %(Tables \ref{tm29blue} and \ref{tm29red}).


\subsection{H$\alpha$ line -- optical depth effects}
\label{sm29haline}

Very noticeably in this nebula is that it shows an exceptionally wide H$\alpha$ profile. Balick (1989) reports a width 
of 11000~km~s$^{-1}$ for this line; Arrieta \& Torres-Peimbert (2003) report a full width (FWZI) of 5000~km~s$^{-1}$ that 
they explain as the result of Raman scattering. The core of the H$\alpha$ line is also anomalous in that it exhibits at 
least two strong components with peaks separated by $\sim 70$~km~s$^{-1}$ (see Figure~\ref{fm29hratios}). 


%%%%% FIGURA 7

\begin{figure}[!t]
\centering
\includegraphics[width=0.99\columnwidth]{storresfig7.eps}
\caption{Evidence that the H$\alpha$ line is optically thick in the nucleus of {M~2-9}. In 
the upper panel the extinction corrected profiles H$\alpha$ and H$\beta$ $\times 3$ lines 
are shown. Zero corresponds to the systemic velocity. In the lower panel the H$\alpha$/H$\beta$ 
ratio of the extinction corrected spectra show excesses in H$\alpha$ at --40 km~s$^{-1}$ and 
+60 km~s$^{-1}$. The uncertainties and the value of the relation for Case B of recombination 
are shown.}
\label{fm29hratios}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% FIGURA 8

\begin{figure}[!t]
\centering
\includegraphics[width=0.97\columnwidth]{storresfig8.eps}
\caption{Dependence on the Balmer intensity ratios on H$\alpha$ optical depth in the nuclear 
region of {M~2-9}. The extinction corrected H$\alpha$/H$\beta$ ratio is = 17.3. It is consistent 
with a value of $\tau$(3s~-~2p) = 8, while H$\gamma$/H$\beta$ = 0.531 with $\tau$(3s~-~2p) = 2. 
In this work a mean value of $\tau$(3s~-~2p) = 5 was adopted.}
\label{fm29optdp}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


The H$\alpha$/H$\beta$ ratio from the extinction corrected spectrum shows two maxima near --40 and +60~km~s$^{-1}$ and 
a minimum at about --10~km~s$^{-1}$. The ratio at minimum is consistent with the predictions of Case B recombination 
(value indicated in the bottom panel of Figure~7). On the other hand, the red shifted peak has a ratio about 6 times 
the value predicted by recombination theory. 

Neither collisional excitation nor extinction can explain the large departure from Case B. Instead, the difference can 
be explained by optical depth effects for transitions to levels n=2. We estimate the optical depth effects by including 
an escape probability
\begin{equation}
P_\mathrm{esc}={1-\exp(-\tau)\over\tau}
\label{em29proba}
\end{equation}
on the Balmer lines. Here 
\begin{eqnarray}
\tau_\mathrm{ul} = \mathrm{constant} \times \lambda f_\mathrm{lu} \Big(n_\mathrm{l}-{n_\mathrm{u}g_\mathrm{l}\over g_\mathrm{u}}\Big) \nonumber \\
\approx \mathrm{constant} \times \lambda f_\mathrm{lu} n_\mathrm{l}\,,~~~~~~~~~~~~
\label{em29optdepth}
\end{eqnarray}
where, $f_\mathrm{lu}$ is the oscillator strength, $n_\mathrm{l}$ and $n_\mathrm{u}$ are the populations of the levels involved 
and $g_\mathrm{l}$ and $g_\mathrm{u}$ the statistical weights. Then, we compute the H~I recombination spectrum using the model 
of Bautista et al. (1998) assuming that optical depths for all the permitted transitions to level n=1 have an infinitely 
large value (Case B approximation) and normalizing the optical depth of the Balmer series to that of the 3s~-~3p transition, 
which is kept as a free parameter. Figure~\ref{fm29optdp} shows the value of the H$\alpha$/H$\beta$ ratio as a function of 
$\tau$(3s~-~2p). It can be seen that the H$\alpha$/H$\beta$ ratio increases monotonically by nearly an order of magnitude as 
$\tau$(3s~-~3p) grows from 0 to 10. The figure also shows that the H$\gamma$/H$\beta$ ratio increases with optical depth too, 
but by a much smaller factor. This explains that traces of the blue shifted components and the red shifted components are 
present in H$\beta$. These components, however, are much smaller and closer to the center of the line that in the case of H$\alpha$. 


%%%%% FIGURA 9

\begin{figure}[!t]
\centering
\includegraphics[width=0.99\columnwidth]{storresfig9.eps}
\caption{Comparison of the Balmer series and [O~III] profiles in the nuclear region 
of {M~2-9}. Both diagrams are on the same velocity scale and are normalized to the 
maximum of the bright line; on the upper panel H$\beta$ and H$\gamma$ have been multiplied 
by 6. The higher [O~III] density corresponds to the central H$\alpha$ minimum.}
\label{fm29profoiii}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Furthermore, the excess in H$\alpha$ line intensity relative to the one predicted for case B can be explained by high 
optical depths for the transitions of level n=2. Therefore, it is proposed that the different observed components in 
these lines are the result of a single optically thick line.

Considering the total flux in the H$\alpha$, H$\beta$ and H$\gamma$ lines it is found that H$\alpha$/H$\beta$ = 17.3 
and H$\gamma$/H$\beta$ = 0.531 after extinction correction. These ratios can be reproduced by two models of recombination 
with approximate $\tau$(3s~-~2p) values of 8.0 and 2.0 respectively. A better approximation in all the H~I spectrum could 
require a more detailed treatment of radiation transfer of the lines. In this work we adopt a mean value of $\tau$(3s~-~2p) = 5.

The optical depth in the spectrum of H~I has an important effect over the emissivities of the predicted lines and therefore 
also for the derived H$^+$ abundance. For example, for $\tau$(3s~-~2p) = 5 the emissivity of the H$\beta$ line is reduced by 
a factor of 2.6 with respect to Case B recombination, then the calculated H$^+$ abundance increases by this same factor. 

In brief, an extinction law in the nucleus characterized by $R_\mathrm{V}=5.0 \pm 0.5$ and $A_\mathrm{V}=2.7 \pm 0.1$ is used 
to correct the nuclear spectrum for extinction. This extinction law corresponds to larger dust grains than those of the typical 
interstellar medium. The exceptionally high value of the intensity of the H$\alpha$ line relative to the predicted recombination 
by Case B is explained by a Balmer line optical depth $\tau$(3s~-~2p) $\sim 5$. Further, we propose that the double peaked 
profile of these lines is also due to optical depth effects.


\subsection{Forbidden line profiles}
\label{sm29forbid}

Balick (1989) found that the central region [N II] lines are well characterized by a bright central Gaussian at the 
systemic velocity, plus two weaker ``satellite'' lines at about $\pm 52$~km~s$^{-1}$. The bright [N II] emission seen 
towards the nucleus is ``circumstellar''. This conclusion is also supported by the very large ratio [N II] 5755/6583 
intensities observed towards the nucleus. The circumstellar [N II] lines are spatially unresolved when the slit is 
oriented E-W (perpendicular to the lobes). Yet, with the slit along the lobes, the two satellite [N II] lines are 
offset about 0.7\arcsec\ from the central star. Thus, the satellite lines probably arise along the nebular symmetry 
axis, and are unlikely to be directly associated to the 2~$\mu$m E-W disk surrounding the star observed by Aspin 
et al. (1988). 

Solf (2000) found that [N II] and [O III] in the nuclear region show 3 components of different velocity, the central 
one at the systemic velocity (low velocity component) and the blue and red shifted components at about $\pm 50$~km~s$^{-1}$ 
(high velocity components), all of which could be resolved spatially along the major axis of the nebula. 

We analyze the complex O$^{++}$ emission line profiles in our spectra. While the 5007 and 4959~\AA\ lines present double 
peaked profiles, with a weak component centered at --17~km~s$^{-1}$ and a bright one centered at --45~km~s$^{-1}$, the 
4363~\AA\ line shows a single component centered at --17~km~s$^{-1}$. The H~I and [O~III] profiles are different; whereas 
for hydrogen the red component is brighter that the blue one, in contrast for oxygen the blue component is stronger than 
the red one. In Figure~\ref{fm29profoiii} we compare the Balmer and [O~III] lines on the same velocity scale. The 
($\lambda\lambda$5007+4959)/$\lambda$4363 ratios in the two zones indicate that there is a region with a very high density 
at --17~km~s$^{-1}$ and another of lower density at --45~km~s$^{-1}$. Due to the velocity displacement between the components, 
it is expected that the higher density component is associated to the region where all the single lines of the spectrum 
are produced. On the other hand, the blue component is formed in a different region.


%%%%% FIGURA 10

\begin{figure}[!t]
\centering
\includegraphics[width=0.99\columnwidth]{storresfig10.eps}
\caption{Observed intensity of the permitted and forbidden Fe~II lines vs. the predicted 
intensities in {M~2-9}. The predicted lines correspond to a $W=10^{-9}$, $T_\mathrm{e}=10^4$~K 
and $N_\mathrm{e}=10^5$~cm$^{-3}$ model. The upper panel corresponds to the lines in the blue 
spectrum and the lower panel to those of the red spectrum.}
\label{fm29fe}
\vspace{0.3cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\subsection{Physical Conditions and kinematics of the emitting region}
\label{sm29phys}

It has been noted that the low value of [O III] 5007/4363 indicates high densities {$\sim 10^{6-7}$~cm$^{-3}$} (Allen \& Swings 
1972; Goodrich 1991). However, a low density region also contributes to the core; the presence of 
the [S II] 6717,6731 lines indicates Ne~$< 10^4$~cm$^{-3}$. There is no significant [N I] 5200 emission in the core; the 
critical density of this line is $\sim 2000$~cm$^{-3}$. 

In our spectra the [N~II], [O~II], [O~III] and [S~II] line ratios were used to determine the temperature and density in the 
nucleus; for this purpose we applied the NEBULAR tasks of the IRAF data reduction package. In addition, we used models of 
Bautista \& Pradhan (1998) for Fe~II, Fe~III, and Fe~IV, and the model of Bautista (2004) for Ni~II. 

As it has been mentioned in \S~\ref{sm29forbid} the [O III] lines also show two components: one centered at {--17~km~s$^{-1}$} 
and another at --45~km~s$^{-1}$ (see Figure~\ref{fm29profoiii}). From the ($\lambda\lambda$5007+4959)/$\lambda$4363 line ratio 
it is found that these components originate in regions of different density. The component at --17~km~s$^{-1}$ corresponds to 
a density $\sim 10^7$~cm$^{-3}$ (that we associate to the nuclear emission proper) while the one at --45 km~s$^{-1}$ to 
$\sim 10^6$~cm$^{-3}$ (that we associate to the micro jets emerging from the nuclear region identified by Balick 1999).

The [Fe~II] and [Ni~II] lines require special attention. The fact that there are permitted Fe~II lines present in the spectrum 
indicates that this ion is undergoing fluorescence excitation (Bautista \& Pradhan 1998). Fluorescence is also suggested by 
the [Ni~II] line intensity ratios, which cannot be reproduced by collisional excitation alone. For the analysis of the [Fe~II] 
lines, those that are less sensitive to fluorescence are selected first; they are usually on the red and infrared part of the 
spectrum (Bautista \& Pradhan 1998). These lines indicate an electronic density of $\sim 10^5$~cm$^{-3}$. To model the lines 
affected by fluorescence, the radiation continuum from 0 to 13.6~eV is approximated by a 40,000 K black body spectrum. The 
intensity of the radiation in the region of the emission line is proportional to the dilution factor
\begin{equation}
W={1\over 4}\Big({R\over r}\Big)^2\,,
\label{em29fdilution}
\end{equation}
where $R$ is the radius of the ionizing star and $r$ is the distance from the emission region to the star. The value of $W$ 
is adjusted to find the best agreement between predicted relative line intensities and the observed fluxes of Fe~II and Ni~II 
lines. It must be noted that variations in the adopted temperature for the black body can be compensated for the most part 
by different values of $W$, such that one gets essentially identical fits to the observed spectrum and predicted line emissivities. 

In Figure~\ref{fm29fe} we plot observed Fe~II permitted and forbidden intensities against the intensities predicted by the 
model with $W=10^{-9}$, $T_\mathrm{e}=10^4$~K and $N_\mathrm{e}=10^5$~cm$^{-3}$. The upper panel corresponds to the blue spectrum, 
the lower to the red spectrum. The agreement with the red region is very good while the agreement with the blue region is only 
moderately satisfactory, although acceptable, given the lower precision of the weak lines and the uncertainties in the radiation 
field that affects the Fe~II lines. However, it should be noted that both forbidden and permitted lines are reasonably well explained 
with the same model.


%%%%%%% TABLE 3

\setcounter{table}{2}
\begin{table}[!t]
\vspace{0.1cm}
\centering
\caption{Comparison of observed and theoretical [Fe~ III] intensities  normalized to $\lambda5085$($a~^3P_1-a~^5D_0$)} 
\label{tm29fe3}
\small
\begin{tabular}{cccc}
\toprule
Wavelength & Transition  & Obs.  & Theoretical  \\
           &             & \multicolumn{2}{c}{I($\lambda$)/I($\lambda$5085)} \\
\midrule 
4607 & $a~^3F_3-a~^5D_4$ &  ~2.87 & ~1.84 \\
4658 & $a~^3F_4-a~^5D_4$ &  17.2~ & 18.3~ \\
4702 & $a~^3F_3-a~^5D_3$ &  14.5~ & ~9.85 \\
4734 & $a~^3F_2-a~^5D_2$ &  ~2.95 & ~4.05 \\
4755 & $a~^3F_4-a~^5D_3$ &  ~2.91 & ~3.43 \\
4770 & $a~^3F_3-a~^5D_2$ &  ~2.38 & ~3.40 \\
4778 & $a~^3F_2-a~^5D_1$ &  ~1.40 & ~1.96 \\
4881 & $a~^3H_4-a~^5D_4$ &  ~2.11 & ~2.15 \\
4931 & $a~^3P_0-a~^5D_1$ &  ~2.96 & ~2.26 \\
5011 & $a~^3P_1-a~^5D_2$ &  ~7.82 & ~5.83 \\
5085 & $a~^3P_1-a~^5D_0$ &  ~1~~  & ~1~~  \\
5270 & $a~^3P_2-a~^5D_3$ &  18.5~ & 11.5~ \\
\bottomrule 
\end{tabular}
\vspace{0.15cm}
\end{table}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%%%% TABLE 4

\setcounter{table}{3}
\begin{table}[!b]
\vspace{0.2cm}
\centering 
\caption{Comparison of   observed and theoretical  [Fe~ IV] intensities normalized to $\lambda5234$ ($a~^4F_{3/2}-a~^4G_{7/2}$)} 
\label{tm29feiv}
\tablecols{4}
\setlength{\tabcolsep}{0.8\tabcolsep}
\footnotesize
\begin{tabular}{cccc}
\toprule
Wavelength  & Transition & Obs. & Theoretical \\
            &            & \multicolumn{2}{c}{I($\lambda$)/ I($\lambda5233$)} \\
\midrule
~4868 & $a~^2F_{5/2}-a~^4G_{9/2,11/2}$ & 0.70 & 0.72~ \\
~4900 & $a~^4F_{9/2}-a~^4G_{7/2}$~~   & 0.57 & 0.56~ \\
\multirow{2}{*}{\Big\{4906}~ & $a~^4F_{9/2}-a~^4G_{11/2}$~   & \multirow{2}{*}{2.40} & \multirow{2}{*}{~1.72\Big\}} \\
     & $a~^4F_{7/2}-a~^4G_{5/2}$~~   &             \\
~5033 & $a~^4F_{3/2}-a~^4G_{7/2}$~~   & 1~~  & 1~~~  \\
~5034 & $a~^4F_{5/2}-a~^4G_{11/2}$~   & 0.75 & 0.74~ \\
\bottomrule 
\end{tabular}
%\vspace{0.2cm}
\end{table}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Our Fe~III model reproduces reasonably well the twelve identified [Fe~III] lines in the blue spectrum. The agreement between 
the model predictions and the observations can be seen in Table~3. %Table \ref{tm29fe3}. 
Here the mean of the dispersion between the ratio of Fe$^{++}$/H$^+$ abundance calculated for each line is less that 40\%, which 
is in accordance to estimated uncertainties in the intensity of these weak lines.

The [Fe~IV] lines in the optical spectrum of {M~2-9} are consistent with the presence of a high density region. The effective 
collision strengths for optical [Fe~IV] lines are very small; thus, high electron densities are required for those lines to 
be visible. Table~4 shows the comparison of model predictions with observations. The agreement is excellent for all the lines 
excepting that at 4906~\AA. This disagreement can be due to blends.


%%%%%% TABLE 5

\setcounter{table}{4}
\begin{table}[!t]
\vspace{0.1cm}
\centering 
\caption{Radial velocity and physical conditions of different ions from forbidden lines in {M~2-9}} 
\label{tm29kin}
\setlength{\tabnotewidth}{0.99\columnwidth}
\tablecols{6}
\setlength{\tabcolsep}{0.95\tabcolsep}
\footnotesize
\begin{tabular}{lccccc}
\toprule
Ion & \multicolumn{2}{c}{Radial}                  & log $N_\mathrm{e}$  & $T$ & Ionization  \\
    & \multicolumn{2}{c}{velocity\tabnotemark{a}} & cm$^{-3}$          & K  & potential \\
 \midrule          
$[$N I$]$    & --11.$\pm$~4. & ~(2)  & -   & -      &  ~0.~~     \\
$[$N II$]$   & ~--6.$\pm$~2. & ~(3)  & 5.6 & ~9500  &  14.53     \\  
$[$N IV$]$   & --17.$\pm$-~~ & ~(1)  & -   & -      &  47.45     \\
$[$O I$]$    & ~--2.$\pm$~0. & ~(2)  & -   & -      &  ~0.~~     \\
$[$O II$]$   & ~--3.$\pm$11. & ~(4)  & 5.6 & ~9500  &  13.62     \\
$[$O III$]$  & --17.$\pm$-~~ & ~(1)  & 7.0 & 11500  &  35.12     \\
$[$Ne III$]$ & --48.$\pm$~0. & ~(2)  & -   & -      &  40.96     \\
$[$S II$]$   & ~--5.$\pm$~9. & ~(4)  & 5.0 & ~9000  &  ~9.98     \\  
$[$Cl II$]$  & ~+0.$\pm$-~~~ & ~(1)  & -   & -      &  12.97     \\ 
$[$Cl III$]$ & ~--3.$\pm$-~~ & ~(1)  & -   & -      &  23.81     \\ 
$[$Ar III$]$ & --15.$\pm$~9. & ~(3)  & -   & -      &  27.63     \\
$[$Cr II$]$  & --18.$\pm$21. & ~(4)  & -   & -      &  ~7.34     \\
$[$Fe I$]$   & --15.$\pm$11. & ~(3)  & -   & -      &  ~0.~~     \\
$[$Fe II$]$  & ~--1.$\pm$~8. & (76)  & 5.0 & ~9000  &  ~7.87     \\
$[$Fe III$]$ & ~--8.$\pm$~6. & (14)  & 6.0 & -      &  16.16     \\
$[$Fe IV$]$  & --18.$\pm$11. & ~(6)  & 6.6 & 11500  &  30.65     \\
$[$Ni II$]$  & ~--0.$\pm$~8. & ~(5)  & -   & -      &  ~7.64     \\
$[$Ni III$]$ & --10.$\pm$10. & ~(2)  & -   & -      &  18.17     \\
$[$Ni IV$]$  & --13.$\pm$-~~ & ~(1)  & -   & -      &  35.17     \\
\bottomrule                                
\noalign{\smallskip}
\tabnotetext{a}{Radial velocity and standard deviation. The number of measured lines is shown in parenthesis.}
\end{tabular}
\vspace{0.2cm}
\end{table}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% FIGURA 11

\begin{figure*}[!t]
\centering
\includegraphics[width=12cm]{storresfig11.eps}
\caption{Behavior of physical conditions relative to the observed radial velocity in 
{M~2-9}.  On the upper left panel the velocity dependence on ionization potential is 
shown for those forbidden species for which more than one line was measured. On the lower 
left panel the electron temperature is presented; on the upper right panel, the electron 
density; and on the lower right panel, the electron pressure. The velocities are given 
relative to the systemic velocity. Solid lines indicate the best linear fit to the data.}
\label{fm29vel}
\vspace{0.32cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


We have determined the mean velocity for the different forbidden ions observed. This determination was carried out from those 
lines that are not ostensibly blended. The observed values are presented in Table~5, %Table\ref{tm29kin},
where we also include the standard deviation and the number of lines available. In the nuclear spectrum practically all lines 
have negative velocities relative to the systemic velocity; they lie in the --48 to +0.2~km~s$^{-1}$ range. The absence of 
positive velocity components is an indication of obscuration around the nucleus. The existence of a dense equatorial disk which 
is obscuring the inner portions of the receding matter has been suggested by Solf (2000). Thus, we can only see the front 
section of the circumstellar region. 

It is also found that there is a relationship between the ionization potential and the radial velocity. This relation is 
presented in Table~5 and Figure~\ref{fm29vel} where only those species with more than one observed (unblended) line are 
displayed. That is, it was found that the velocities for different lines of the same ion are comparable, and that these 
velocities show a correlation with the ionization potential, where the more negative velocities correspond to higher 
potentials, and the positive ones to the neutral species. Furthermore, the emission lines originate in regions of a wide 
range of densities, as presented in this same table. The electron density varies from about $10^5$~cm$^{-3}$ in the region 
of low ionization ([S II] emitting region) to $10^7$~cm$^{-3}$ in the [O~III] emitting region for temperatures from 9000 to 
11500~K, respectively. These values are also included in Table~5. This result is in agreement with the Calvet \& Cohen (1978) 
estimates of $N_\mathrm{e}$ and $T_\mathrm{e}$ determined from the [OIII], [NII] and [OI] lines in the nucleus. They further 
argue the existence of ionization stratification, where a dense [OIII] zone is surrounded by a less dense [OI] zone. That is, 
there is also a correlation of physical conditions with radial velocity; where the lines of higher ionization species are 
emitted in regions of much higher density than those of the neutral species. This dependence is displayed in Figure~\ref{fm29vel}, 
where we also show evidence of temperature and electron pressure correlation. This well defined negative gradient of ionization 
potential, electron density, electron temperature and gas pressure with radial velocity was already shown in a preliminary 
work by Arrieta \& Torres-Peimbert (2000). 

Although the spectral data reported here refer to a $2\arcsec \times 10\arcsec$ window, the observed line emission 
is produced only in the nuclear region, which is significantly brighter than the surrounding region; that is, we are 
observing a circumstellar envelope. We adopt a size of the observed circumstellar envelope as that of the [S~II] 
($\sim 1.5\arcsec$ diameter) equatorial region imaged with HST (Figure~\ref{fm29image}), which at a distance of 650 pc 
corresponds to a circumstellar structure of $\sim 490$ AU in radius. The fast moving material (high density, and high 
ionization), at a mean velocity of $\sim 30$~km~s$^{-1}$, can traverse the ionized circumstellar region in $\sim 80$ years. 

We try to describe in a very simple manner the geometry of the optical emitting circumstellar region. Comsidering that 
we observe only negative velocities we rule out the possibility of looking at a complete shell; next we consider the 
case of the front half of a spherical shell, in this case the expansion (deprojected) velocities can be 3 times larger 
than the observed ones; moreover, the profiles would indicate a significant contribution at the systemic velocity. 
Therefore we find it more likely that we are observing a toroidal geometry, an extension of the molecular torus that is 
present very close to the star, and extends to large enough distances to pinch the waist of the lobes.

We place the observed high ionization species ([O~III], [Ne~III], etc.) in the inner part of the torus that is being 
shocked by the fast wind, and the less ionized species ([S~II], [O~I], etc.) in the external part, expanding at a lower 
velocity, consistent with the velocity of the CO torus ($\sim 7$~km~s$^{-1}$). That is, the optical emission lines are 
produced in the inner regions of the torus that is surrounding the central star. This picture is consistent with a very 
extended torus where the inner region is producing the optical emission lines and the radio continuum emission (Kwok 
et al. 1985), followed by the warm H$_2$ component (observed by Smith et al. 2005), that continues to a CO torus (observed 
by Zweigle et al. 1997), altogether forming a warm dusty torus at 260~K (Smith \& Gehrz 2005).

For the case of a quasi-steady state structure, we can estimate the width of the ionized region of the circumstellar 
configuration. From conservation of mass considerations the relative sizes of the emitting regions can be derived 
\begin{equation}
{r_2^2 v_2 n_2} \propto {r_1^2 v_1 n_1}\,,
\label{em29massloss}
\end{equation}
\noindent where $r_1$ and $r_2$ refer to the internal and external radii of the emitting region. From the values reported 
in Table~5 and the linear fit presented in Figure~\ref{fm29vel} we can adopt $v_1 \sim 30$~km~s$^{-1}$ and $v_2 \sim 5$~km~s$^{-1}$ 
as the inner and outer expanding velocities and $n_1 \sim 10^{6.7}$ and $n_2 \sim 10^5$ as the densities. We derive 
$r_2 /r_1 \sim 4.5$. Therefore an estimate of the size of the nuclear emitting region can be determined, which corresponds 
to $r_2 \approx 490$ AU and therefore to $r_1 \approx 110$ AU. 

On the other hand, the small scale of the emitting region ($\sim 490$ AU, allows us to propose that observable changes 
in the kinematical behavior of the optical forbidden lines might take place in timescales of the order of $\sim 40$ years.


%%%%%% TABLE 6

\setcounter{table}{5}
\begin{table}[!t]
\vspace{0.1cm}
\centering 
\caption{Abundances in the {M~2-9} nucleus, and comparison to solar abundances}
\label{tm29abun}
\setlength{\tabnotewidth}{0.99\columnwidth}
\tablecols{7}
\setlength{\tabcolsep}{0.64\tabcolsep}
\scriptsize
\begin{tabular}{lcccccc}
\toprule
X & X$^0$/H$^+$ & X$^+$/H$^+$ & X$^{++}$/H$^+$ & X$^{3+}$/H$^+$ & X/H & X/H  \\
  &             &            &               &               &     & solar\tabnotemark{a} \\
\midrule 
He & ...      & 1.1(--1)  &           &           & 1.1(--1)  & 8.5(--2) \\
N  & 1.0(--5) & 4.1(--5)  & 2.9(--5)  & ...       & 8.0(--5)  & 8.3(--5) \\
O  & 2.5(--6) & 2.2(--4)  & 1.3(--4)  & ...       & 3.5(--4)  & 6.7(--4) \\
Ne & ...      & ...       & 2.8(--7)  & ...       & ...       &          \\
S  & ...      & 5.9(--6)  & 5.8(--6)  & ...       & 1.2(--5)  & 2.1(--5) \\
Ar & ...      & ...       & 7.7(--7)  & ...       &           &          \\
Fe & ...      & 5.4(--7)  & 6.1(--6)  & 2.6(--6)  & 9.2(--6)  & 3.2(--5) \\
Ni & ...      & 1.8(--8)  & 4.2(--7)  & 1.4(--7)  & 5.8(--7)  &          \\
\bottomrule 
\noalign{\smallskip}
\tabnotetext{a}{From Grevesse \& Sauval (1998).}
\end{tabular}
\vspace{0.4cm}
\end{table}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\subsection{Abundances} 
\label{sm29abund}

The estimate of the nebular ion abundances is generally carried out by adopting a mean value for density and temperature 
and then calculating the line emissivities for all species. This approximation is reasonable in most cases where all 
diagnostics for each ion agree within their uncertainty. By contrast, in {M~2-9} we find a clear density gradient with 
distance from the central star. Thus, the ionic abundances were calculated here by adopting a varying density through 
each emission region while the temperature is kept constant at $10^4$~K. In this approximation the density diagnostic 
has been used for each available ion. The calculated abundances for each ion are listed in Table~6. %Table \ref{tm29abun}. 


%%%%% FIGURA 12

\begin{figure*}[!t]
\centering
\includegraphics[width=14.6cm]{storresfig12.eps}
\caption{Image of the planetary nebula {M~1-91} (from Manchado et al. 1996). The size 
is $45 \times 4.6\arcsec$. Our slit covered the central 2\arcsec\ along the major axis.} 
\label{fm191image}
\vspace{0.3cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


The abundances of O$^0$ and N$^0$ were computed including the contributions of continuum fluorescence from the model 
of Bautista (1999). To compute the abundances of Ni~III and Ni~IV we adopted the models of Bautista (2001) and Mel\'endez 
\& Bautista (2005), respectively.

The comparison with solar abundances yields $\mathrm{N/O_{M~2-9}} > \mathrm{N/O}_\odot$; this value is common for this type 
of objects since the alpha mechanism is transforming C into N. Most of the abundances presented in %Table \ref{tm29abun}
Table~6 are in reasonable agreement with solar abundance values, with the exception of Fe and Ni. The low abundance of 
these two elements can be explained by depletion into dust (Georgiev at al. 2006).


\section{M~1-91} 
\label{sm191}

The extreme bipolar planetary nebula \mbox{M~1-91} (also known as {M~4-15}, {He 2-437} and {PN~G061.3+ 03.6}) is morphologically very 
similar to \mbox{M~2-9}. In fact is is even more extreme in its aspect ratio (3.6:1). It is among the most collimated PNe known. 
Although very similar to {M~2-9}, its surface brightness is lower and thus has been less extensively studied. A striking 
difference is that the knots in {M~1-91} are point-symmetric, not plane-symmetric like those near the center in {M~2-9} 
(Figure~\ref{fm191image}).

From deep images it is found that it has a size of $45 \times 4.6\arcsec$ (Machado et al. 1996). Carsenty \& Solf (1983) 
estimate that the inclination angle between the bipolar axis and the plane of the sky is about 15$^{\circ}$. {M~1-91} is not 
only morphologically very similar to {M~2-9}, but its nuclear spectrum is also quite analogous. It is smaller in size and 
considerably fainter. Guerrero et al. (2000) detect H$_2$ only marginally in the lobes of the nebula.

Trammell et al. (1996) report that the lobes are significantly polarized in H$\alpha$. However the high [O~III] 
$\lambda\lambda$ 4959,5007/4363 ratio ($\sim 89$ is the eastern lobe, $\sim 67$ in the western lobe) compared to the one 
found in the nucleus ($\sim 8$), indicates that part of the [O~III] emission in the lobes is intrinsic and not due to dust 
scattered from the nucleus (Goodrich 1991).

Calvet \& Cohen (1978) estimate a distance of 3.1~kpc for this object, while Cahn \& Kaler (1971) and Maciel (1984) estimate 
a distance to the planetary nebula {M~1-91}, of 7.79 and 7.0 kpc, respectively; 
the methods used to determine these values are valid only for optically thin planetary nebulae. As will be discussed later 
we can conclude that the nucleus of {M~1-91} is optically thick, which is another characteristic in common with \mbox{M~2-9}. Rodr\'iguez, 
Corradi, \& Mampaso (2001) adopt a distance for {M~1-91} of 7.8 kpc that implies a height above the galactic plane of $z=440-490$ pc, which 
is still rather high compared to the mean distance to the galactic plane of bipolar planetary nebulae of $<$$z$$>$ = 130~pc 
(Corradi \& Schwartz 1995); therefore, it is possible that the distance to {M~1-91} has been overestimated. There are two ways to 
determine an approximate distance to this object: (a) by assuming that it has the same physical size as {M~2-9}, in which case it 
would be 2.7 times farther away than {M~2-9}, that is at 1.7~kpc, and (b) by assuming that its height above the galactic 
plane is 130~pc (the mean value for bipolar PNe), in which case this would correspond to 2.1~kpc. We will adopt a distance 
of 1.9~kpc for this object.


\subsection{Nucleus of M~1-91}

The spectrum of the core of {M~1-91} closely resembles that of the {M~2-9} core (Calvet \& Cohen 1978) and H$\alpha$ emission 
line also shows extended wings (Goodrich 1991). Also like {M~2-9}, it shows weak [S II] 6717, 6731 and a high  [O III] 4363/5007 
ratio. The lines of [Fe II] are also seen in {M~1-91}, although they are relatively weaker than in {M~2-9}, with somewhat weaker 
[O III] and iron lines, and stronger [O~I] and He I lines.

Rodr\'iguez et al. (2001) estimated densities in the nucleus of {M~1-91} to be in the range of 10$^{3.3}$ to 10$^{6.3}$~cm$^{-3}$ 
depending on the ion under consideration. They propose that the range in densities indicates that some phenomenon of mass 
loss or mass motion is taking place in the nucleus.

In contrast to {M~2-9} and of most young PNe, the nucleus of {M~1-91} does not show a molecular component. Josselin et al. (2000) 
do not detect CO. The absence of an intense molecular component could be due to a hotter central star in {M~1-91}.

The radio continuum images only show a compact core, which coincides with the optical stellar core (Lee, Lim, \& Kwok 2007). Along 
the axis they resolve the source to be from 34 to 77 mas, depending on the frequency. The spectral index is of 0.7 which is 
compatible with an isothermal stellar wind with a constant expansion velocity. Lee et al. (2007) for an isotropic mass loss, derive 
a mass loss rate of $\dot M \sim 9 \times 10^{-5}~M_{\odot}$~yr$^{-1}$. However, they revise this value by proposing that this 
source is elongated possibly by collimated winds; in this case it corresponds to $\dot M \sim 10^{-6}~M_{\odot}$~yr$^{-1}$. 
The mass loss rate derived is at least an order of magnitude higher than expected, which they take as an indication that the 
object is a symbiotic star. They suggest that the radio core represents a collimated outflow responsible for shaping the bipolar 
morphology of the nebula. 

\vspace{0.1cm}

\subsection{Systemic velocity} 
\label{sm191system}

Our slit was located along the major axis, and covered most of the nebula. Also in this case, no significant continuum 
was measured nor any absorption lines were detected in the spectrum of the nucleus of {M~1-91}; therefore the systemic 
velocity determination was carried out under the assumption that the nucleus has the same velocity as the lobes.


%%%%% FIGURA 13

\begin{figure}[!t]
\centering
\includegraphics[width=0.99\columnwidth]{storresfig13.eps}
\caption{Determination of the systemic velocity of \mbox{M~1-91}. The dashed line joins 
both condensations and from it we derive the velocity at the nucleus. The velocity 
difference is  --26~km~s$^{-1}$ for the systemic velocity, which corresponds to a 
heliocentric velocity of $-10 \pm 5$~km~s$^{-1}$.}
\label{fm191posvel}
\vspace{0.4cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%% FIGURA 14

\begin{figure*}[!t]
\centering
\includegraphics[width=14cm]{storresfig14.eps}
\caption{Blue spectrum of the nucleus of {M~1-91}. The vertical scale has been expanded 
to show the faint lines (note that the expansion is different in each panel). The 
identified lines are marked.}
\label{fm191blue}
\vspace{0.3cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


%%%%%%% TABLE 7

\setcounter{table}{6}
\begin{table*}[!t]
\centering 
\caption[]{Emission lines in the nucleus of {M~1-91}}
\label{tm191blue}
\setlength{\tabnotewidth}{0.99\textwidth}
\tablecols{11}
%\setlength{\tabcolsep}{0.9\tabcolsep}
\scriptsize
\begin{tabular}{cccclcccccl}
\toprule
             &                       & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{a}  & FWHM  &                                    &  &              &                       & I$_\lambda$/I$_\mathrm{H\beta}$\tabnotemark{a}  & FWHM  & \\
$\lambda_0$  & $\lambda_\mathrm{obs}$  & ext corrected                                & \AA   & \multicolumn{1}{c}{Identification} &  & $\lambda_0$  & $\lambda_\mathrm{obs}$  & ext corrected                                & \AA   & \multicolumn{1}{c}{Identification} \\
\cmidrule{1-5}
\cmidrule{7-11} 
3721.94 & 3721.48 & 0.047 & 0.5 & H14            &  & 5158.78 & 5158.26 & 0.027 & 2.0 & [Fe II]19F \\
3726.03 & 3725.68 & 0.193 & 0.6 & [O II]         &  & 5191.80 & \multirow{2}{*}{{\normalsize \{}5190.99} & \multirow{2}{*}{0.012} & \multirow{2}{*}{0.9{\normalsize \}}} & O I bl\\              
3728.81 & 3728.33 & 0.092 & 0.5 & [O II]         &  & 5191.82 &         &       &     & [Ar III]3F bl\\        
3734.37 & 3733.96 & 0.193 & 1.2 & H13            &  & 5233.45 & \multirow{4}{*}{\Bigg\{5232.95} & \multirow{4}{*}{0.012} & \multirow{4}{*}{1.0\Bigg\}} & O II bl\\  
3750.15 & 3749.70 & 0.130 & 1.7 & H12            &  & 5233.76 &         &       &     & [Fe IV]F bl\\      
3797.90 & 3797.37 & 0.095 & 1.2 & H10            &  & 5233.70 &         &       &     & N I bl\\   
3869.04 & 3868.19 & 0.459 & 0.8 & [Ne III]1F     &  & 5234.03 &         &       &     & Ne I bl\\      
3889.05 & 3888.41 & 0.281 & 1.3 & H8             &  & 5261.62 & 5261.12 & 0.021 & 1.1 & [Fe II]19F \\   
3967.79 & 3967.11 & 0.148 & 1.3 & [Ne III]1F     &  & 5270.40 & 5269.92 & 0.092 & 1.1 & [Fe III]   \\     
3970.07 & 3969.56 & 0.193 & 0.9 & H$\epsilon$    &  & 5273.35 & 5272.51 & 0.003 & 1.0 & [Fe II]18F \\  
4101.73 & 4101.20 & 0.311 & 0.9 & H$\delta$      &  & 5333.65 & 5333.21 & 0.013 & 1.6 & [Fe II]19F \\     
4287.39 & 4286.85 & 0.009 & 0.7 & [Fe II]7F      &  & 5363.80 & 5362.91 & 0.009 & 2.1 & O II \\    
4340.46 & 4339.90 & 0.459 & 0.9 & H$\gamma$      &  & 5376.45 & 5375.90 & 0.015 & 1.8 & [Fe II]19F \\     
4359.33 & 4358.79 & 0.011 & 0.7 & [Fe II]7F      &  & 5527.34 & \multirow{2}{*}{{\normalsize \{}5526.81} & \multirow{2}{*}{0.003} & \multirow{2}{*}{1.1{\normalsize \}}} & [Fe II]17F bl\\    
4363.21 & 4362.58 & 0.361 & 0.8 & [O III]3F      &  & 5527.57 &         &   	  &     & Fe II bl\\   
4413.78 & 4413.04 & 0.047 & 1.2 & [Fe II]7F      &  & 5754.64 & 5753.85 & 0.061 & 1.3 & [N II]3F \\   
4416.27 & 4415.79 & 0.006 & 0.7 & [Fe II]6F      &  & 5875.66 & 5874.95 & 0.308 & 1.2 & He I  \\  
4471.50 & 4470.90 & 0.089 & 1.0 & He I           &  & 5931.89 & \multirow{3}{*}{\big\{5931.49} & \multirow{3}{*}{0.008} & \multirow{3}{*}{2.1\big\}} & Si I bl\\
4607.03 & 4606.11 & 0.025 & 1.0 & [Fe III]3F     &  & 5932.35 &         &       &     & N II bl\\   
4658.05 & 4657.53 & 0.132 & 1.1 & [Fe III]3F     &  & 5932.20 &         &       &     & Si I bl\\  
4701.53 & 4700.90 & 0.061 & 1.0 & [Fe III]3F     &  & 6000.55 & \multirow{2}{*}{{\normalsize \{}5999.94} & \multirow{2}{*}{0.006} & \multirow{2}{*}{1.5{\normalsize \}}} & Fe I bl\\   
4713.22 & 4712.56 & 0.027 & 0.7 & He I           &  & 6000.93 &         &       &     & Ne I bl\\  
4754.69 & 4754.26 & 0.030 & 0.9 & [Fe III]3F     &  & 6046.39 & 6045.58 & 0.005 & 0.8 & O I \\ 
4769.43 & 4768.87 & 0.022 & 1.2 & [Fe III]3F     &  & 6300.30 & 6299.80 & 0.068 & 1.3 & [O I]1F \\ 
4814.51 & \multirow{2}{*}{{\normalsize \{}4813.78} & \multirow{2}{*}{0.022} & \multirow{2}{*}{1.2{\normalsize \}}} & O II bl   &  & 6312.10 & 6311.31 & 0.141 & 1.2 & [S III] \\  
4814.52 &         &       &     & [Fe II]20F bl  &  & 6347.11 & 6346.31 & 0.012 & 1.1 & Si II(2) \\   
4861.32 & 4860.89 & 1.000 & 1.0 & H$\beta$       &  & 6363.78 & 6363.17 & 0.021 & 1.2 & [O I]1F \\
4906.00 & 4905.40 & 0.003 & 1.0 & O I            &  & 6371.37 & 6370.66 & 0.008 & 1.4 & Si II(1) \\   
4921.93 & 4921.22 & 0.027 & 1.0 & He I           &  & 6548.04 & 6547.30 & 0.133 & 1.3 & [N II]1F \\ 
4930.53 & 4930.05 & 0.015 & 1.8 & [Fe III]1F     &  & 6562.82 & -       & 7.541 &     & H$\alpha$\tabnotemark{\dag} \\ 
4958.91 & 4958.28 & 0.781 & 1.0 & [O III]1F      &  & 6583.46 & 6582.69 & 0.441 & 1.4 & [N II]1F \\ 
5006.84 & 5006.20 & 2.325 & 1.1 & [O III]1F      &  & 6678.15 & 6677.17 & 0.098 & 1.3 & He I(46) \\  
5011.25 & 5010.66 & 0.044 & 1.1 & [Fe III]1F     &  & 6716.44 & 6715.81 & 0.022 & 1.2 & [S II]2F \\  
5015.68 & 5015.02 & 0.033 & 1.1 & He I           &  & 6730.81 & 6730.09 & 0.031 & 1.3 & [S II]2F \\  
5018.39 & 5017.49 & 0.018 & 1.0 & O II           &  &         &         &       &     &          \\
\bottomrule
\noalign{\smallskip}
\tabnotetext{a}{I(H$\beta$)$_\mathrm{obs}=1.85 \times 10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ and I(H$\beta$)$_\mathrm{extcorr}=4.40 \times 10^{-11}$~erg~cm$^{-2}$~s$^{-1}$.}
\tabnotetext{\dag}{Several components.}
\end{tabular}
\vspace{0.2cm}
\end{table*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Figure~\ref{fm191posvel} shows isocontours on the position-velocity diagram of the long slit echelle spectrum at H$\alpha$. 
Here a very bright central region with a double peaked profile is seen, as well as the eastern and western lobes. The arrows 
indicate the laboratory wavelength of H$\alpha$, ($\lambda _0$) and the position of the central part of the nebula. The 
difference in velocity of the arrows is $\sim -26 \pm 5$~km~s$^{-1}$. Correcting for the earth's motion (+16.5~km~s$^{-1}$) 
it is found that the heliocentric velocity of the {M~1-91} system is $-10 \pm 5$~km~s$^{-1}$ and the 
$V_\mathrm{LSR}=+9 \pm 5$~km~s$^{-1}$. In the following sections all the discussions of velocities are referred to the systemic 
velocity of the nebula.


\subsection{Spectrum description} 
\label{sm191spectrum}

Motivated by the similarity of their spectra the same type of study as performed for {M~2-9} and described in \S~\ref{sm29} 
was carried out for {M~1-91}. Unfortunately it is significantly fainter, and thus it was not possible to derive as much 
information as for \mbox{M~2-9}. The following results are based on our spectral observations of the nuclear zone. Details of 
the observations and reductions have been given in \S~\ref{sobserv}. 

As in the case of {M~2-9}, the nucleus of {M~1-91} shows a faint continuum without absorption lines; however the number 
of emission lines is more limited. The blue spectrum from 4150 to 6850 is presented in Figure~\ref{fm191blue}. The 
emission spectrum shows permitted lines (H~I, He~I, N~I, N~II, O~I, O~II, Ne~I, Si~I, Fe~I and Fe~II) and forbidden 
lines ([N~II], [O~I], [O~II], [O~III], [Ne~III], [Ar~III], [S~II], [S~III], [Fe~II], [Fe~III] and [Fe~IV]). It was 
also useful to resort to the code EMILI (Sharpee et al. 2003) to identify the faint lines. The code was used for a 
set of different temperatures (9000 to 11000) and densities (10$^4$ to 10$^{6.5}$~cm$^{-3}$). The profiles of the Balmer 
lines and the [O~III] 5007, 4959 lines show different unresolved regions. The {M~2-9} and {M~1-91} spectra are similar, 
as far as the signal-to-noise ratio allowed such a comparison. Practically all the {M~2-9} emission lines stronger than 
I$_\lambda$/I$_\mathrm{H\beta} > 0.05 $ are present in {M~1-91}. 
 
In Table~7 the observed spectral lines are listed: Columns~(1) and (2) present the laboratory and observed wavelength 
of the identified lines; Column~(3) shows the extinction corrected intensities relative to H$\beta$ (the details about 
the extinction law are given in \S~\ref{sm191extinc}); Column~(4) lists the FWHI (not corrected for instrumental 
broadening); and in Column~(5) the line identification is given.


\subsection{Balmer and [O~III] profiles} 
\label{sm191profiles}

Analogously to the case of {M~2-9} the Balmer and [O~III] 5007 and 4959 line profiles show several components. In the 
Balmer series it can be noticed that H$\alpha$ shows two well separated components and this separation is less defined 
with increasing upper level. Our {M~1-91} spectrum reaches only wavelengths of 6800~\AA, and therefore it was not possible 
to compare the observed Paschen profiles as was the case for \mbox{M~2-9}. In the case of the [O~III] lines, the two components 
are not so well separated as in the hydrogen case, but it can be appreciated that the [O~III] 5007 and 4959 lines have 
two components, while [O~III] 4363 has a single component. Also, in hydrogen the red component is stronger than the blue 
one, in contrast to what is observed in oxygen. These same characteristics are present in the nucleus of {M~2-9} 
(\S\S~\ref{sm29haline} and \ref{sm29forbid}).


\subsection{Interstellar extinction} 
\label{sm191extinc}

To derive the visual extinction, $A_\mathrm{V}$, the hydrogen emission lines were compared with predictions for Case B 
recombination using a normal interstellar extinction law.

Assuming a mean density and temperature of $10^6$~cm$^{-3}$ and $10^4$~K, respectively, and adopting the extinction curve 
by Cardelli et al. (1989) with $R_\mathrm{V}=3.1$, it is found that while the observed ratio of H$\alpha$ to 
H$\beta$ suggests $A_\mathrm{V} \sim 5$, all the other Balmer lines are consistent with $A_\mathrm{V}=2.4 \pm 0.3$. The low 
signal-to-noise ratio in these lines did not allow us to explore the case of an anomalous extinction curve, as in the case 
of {M~2-9}. Therefore, a value of $A_\mathrm{V}=2.4 \pm 0.3$ was adopted from the H$\gamma$/H$\beta$, H$\delta$/H$\beta$, 
and H$\epsilon$/H$\beta$ ratios. This value was used to correct the spectrum for extinction.

As in {M~2-9}, the H$\alpha$ line excess can be explained as an effect of high optical depth to levels n=2. In 
Figure~\ref{fm191ratiohi} it is shown that the H$\alpha$ line exhibits two strong components with peaks separated by 
$\sim 50$~km~s$^{-1}$. Observing the H$\alpha$/H$\beta$ ratio in this same figure, one finds two maxima of this ratio 
about --30 and +40~km~s$^{-1}$ and a minimum centered on the systemic velocity. This ratio in the minimum is consistent 
with the recombination predictions in Case B (shown in the figure). On the other hand, the highest value of the 
H$\alpha$/H$\beta$ ratio is about 3 times higher than predicted from case B recombination theory.


%%%%% FIGURA 15

\begin{figure}[!t]
\centering
\includegraphics[width=0.99\columnwidth]{storresfig15.eps}
\caption{H$\alpha$ and H$\beta$ $\times 3$ in {M~1-91}. The extinction corrected profiles of H$\alpha$ 
and H$\beta$ $\times 3$ in {M~1-91} are shown in the upper panel. The velocity is given relative to the 
systemic velocity. In the lower panel the extinction corrected H$\alpha$/H$\beta$ ratio is shown. At 
--30~km~s$^{-1}$ and +40~km~s$^{-1}$ excesses of the H$\alpha$/H$\beta$ ratio can be appreciated. The 
uncertainties are marked, and the recombination case B value is also included. This diagram is evidence 
that the H$\alpha$ line is optically thick.}
\label{fm191ratiohi}
\vspace{0.2cm}
\end{figure}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


Comparing the total extinction-corrected H$\alpha$ and H$\beta$ flux it is found that H$\alpha$/H$\beta$ = 7.5. This 
ratio can be reproduced with a case B recombination model with $\tau$(3s~-~2p) $\sim 4$ (as derived from Figure~\ref{fm29optdp}).


%%%%%%% TABLE 8

\setcounter{table}{7}
\begin{table}[!t]
\vspace{0.1cm}
\centering 
\caption{Radial velocity and physical conditions of different ions from forbidden lines in {M~1-91}} 
\label{tm191kin}
\setlength{\tabnotewidth}{0.99\columnwidth}
\tablecols{5}
%\setlength{\tabcolsep}{1.5\tabcolsep}
\small
\begin{tabular}{lcccc}
\toprule
Ion & \multicolumn{2}{c}{Radial velocity\tabnotemark{a}} & log $N_\mathrm{e}$ & Ionization \\
    &                   &                                & cm$^{-3}$          & potential \\
\midrule
$[$N II$]$   &  --11.$\pm$~3.  & ~(3)  & 5.15\tabnotemark{b}  &  14.53   \\
$[$O I$]$    &  ~--0.$\pm$~3.  & ~(2)  & --                   &  ~0.~~   \\
$[$O II$]$   &  ~--7.$\pm$~5.  & ~(2)  & 3.5~~                &  13.62   \\
$[$O III$]$  &  --14.$\pm$~2.  & ~(3)  & 6.4\tabnotemark{b}~  &  35.12   \\
$[$Ne III$]$ &  --33.$\pm$~7.  & ~(2)  & --                   &  40.96   \\
$[$S II$]$   &  ~--4.$\pm$~2.  & ~(2)  & 3.34\tabnotemark{b}  &  10.36   \\
$[$S III$]$  &  --12.$\pm$~-~  & ~(1)  & --                   &  23.33   \\
$[$Fe II$]$  &  --10.$\pm$~8.  & (11)  & 6.0\tabnotemark{b}~  &  ~7.87   \\
$[$Fe III$]$ &  --10.$\pm$10.  & ~(8)  & 6.5\tabnotemark{b}~  &  16.16   \\
$[$Fe IV$]$  &  --20.$\pm$~-~  & ~(1)  & --                   &  30.65   \\
$[$Ar III$]$ &  --22.$\pm$~-~  & ~(1)  & --                   &  40.74   \\
\bottomrule                           
\noalign{\smallskip}
\tabnotetext{a}{Radial velocity and standard deviation. The number of measured lines is shown in parenthesis.}
\tabnotetext{b}{Rodr\'iguez et al. (2001).}
\end{tabular}
\vspace{0.2cm}
\end{table}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


\subsection{Physical conditions and kinematics of the emitting region}
\label{sm191phys}

The forbidden line ratios were used as density diagnostics, applying the NEBULAR tasks of IRAF. The [O~II] and 
[S~II] density determinations were obtained from 3729/3726 and 6717/6731 ratios, respectively. The [Fe~II], 
[Fe~III], [N~II] and [O~III] densities were taken from the work by Rodr\'iguez et al. (2001).

The electron density varies from about 2200~cm$^{-3}$ in the region of low ionization ([O~II] emitting region) to 
$3 \times 10^6$~cm$^{-3}$ in the [O~III] emitting region (Table~8). %(Table \ref{tm191kin}

All of the forbidden lines show negative velocities, which are indicative of an obscuring disk around the core. 
The estimate of the velocity for different ions where more than one line is observed is in the $-33$ to 0~km~s$^{-1}$ 
range relative to the systemic velocity. These velocities show a correlation with the ionization potential, where 
the more negative velocities correspond to higher potentials, and the positive ones to the neutral species.


%%%%% FIGURA 16

\begin{figure*}[!t]
\centering
\includegraphics[width=13.8cm]{storresfig16.eps}
\caption{Behavior of the observed radial velocity with ionization potential and density 
for M1-91. On the left panel is presented the kinematical behavior of the observed radial 
velocity with ionization potential in the observed forbidden lines, for which more than 
one line was measured. On the right panel the electron density behavior with observed 
radial velocities is displayed. The velocities are given relative to the systemic velocity. 
The solid line indicates the best fit to the data.}
\label{fm191vel}
\vspace{0.3cm}
\end{figure*}

%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%


There is a correlation between density and velocity in the nuclear spectrum of this object. Again similarly to 
{M~2-9}, we find that the behavior of density is not only correlated to the ionization potential but also to the 
observed radial velocity in each ion. In Figure~\ref{fm191vel} the relative velocity is plotted for the different 
species for which we observe more than one line as a function of ionization potential and electron density. 

For this object, we also propose that the observed optical emission lines are produced on the front part of the 
object, while the receding part is obscured by the dense torus. We can only speculate that the ionized region is 
the inner part of the torus irradiated by the star. We propose that the high density emitting zones are the regions 
closer to the nucleus, while the neutral or those of lower degree of ionization are at larger distances from the star. 

The size of the core emitting region is difficult to estimate. The 1.3 cm radio continuum observations yield a size 
of 0.077\arcsec\ (Lee et al. 2007). However, we can conservatively assume that the total diamter of the ionized 
circumstellar region is 10 times this value, that is, of $\sim 0.8\arcsec$. This size is compatible with the $\sim 2\arcsec$ 
size of the direct images available, which are upper limits to the apparent size. At a distance of 1.9~kpc, $0.8\arcsec$ 
corresponds to a radius $r_2 \approx 760$ AU. At the mean velocity of 33~km~s$^{-1}$ the distance from the star 
to the outskirts of the ionized circumstellar torus can be crossed in 115 years. Then it is conceivable that detectable 
changes in the kinematical behavior of the system can be expected in a fraction of this time. We suggest that in 
$\sim 30$ years there might be significant changes in the velocity distribution of the emitting region. 

Alternatively, in the case of quasi-equilibrium, where we assume that constant mass loss is occuring and deceleration 
of the wind takes place, we can assume that the equation~\ref{em29massloss} is valid. For the high and low velocities 
$v_1 \sim 33$~km~s$^{-1}$, and $v_2 \sim 5$~km~s$^{-1}$, and for $n_1 \sim 10^{6.5}$~cm$^{-3}$, and $n_2 \sim 10^{3.3}$~cm$^{-3}$, 
we can derive $r_2/r_1 \sim 32$. For $r_2 \approx 760$ AU (low ionization region) we obtain $r_1 \approx 50$ AU (high 
ionization region), which would be the inner region where the fast wind is impinging on the dense torus.


\section{Summary and conclusions}

For M2-9 from the hydrogen Balmer and Paschen series in the nucleus, it was found that the extinction curve has 
unusual characteristics, different from those of a gas under the action of a strong radiation field. This curve is 
described by $R_\mathrm{V}=5.0$, $A_\mathrm{V}=2.7$ and explains the discrepancies reported in the literature. The 
double peak found in the profiles of the Balmer lines had been previously explained by accretion discs like those 
found in symbiotic stars. According to those models the localized condensations in the disc produces the bright 
peak in the profile. As the hot spot rotates about of the star, the bright peak in the profile is blue and red 
shifted. In this work, the possibility that a fast moving hot spot in an accretion disc produces the asymmetry 
found in these profiles is not favored; this is based on observations of this object carried out at different 
epochs, only for the purpose of monitoring the profile morphology. It was found that in all spectra (including 
the one reported by Balick 1989) the bright peak is always red shifted and the intensity ratio between both peaks 
remains unchanged. In this section it was found that the double profile can be associated to optical depth effects. 
From the H$\alpha$/H$\beta$ and H$\gamma$/H$\beta$ line ratios an optical depth of $\tau_\mathrm{H\alpha} \sim 5$ in 
the line H$\alpha$ was estimated. In addition, from the (5007+4959)/4363 [O~III] line ratio it was found that there 
are two different zones in the nuclear regions not spatially resolved, a component of very high density and velocity 
--20~km~s$^{-1}$ and another less dense with velocity of --45~km~s$^{-1}$. Here it is proposed that the denser component 
is associated to the circumstellar torus, while the less dense component corresponds to the collimated structure 
seen in the [O~III]~5007 images obtained with HST by Balick et al. (1997). The forbidden emission 
lines in the core present negative velocities (including [O~III] that has a complex profile); that is, we are only 
observing the approaching side of the core, and presumably there is obscuring material that does not allow us to 
see the receding part. From the [S~II] and [N~II] narrow filter HST observations it was found that in these ions the 
region has an extended morphology in the equatorial direction. It is here proposed that this extended region is the 
one that emits with these peculiar characteristics. The upper limit to the size of the ionized region is that of the 
[S II] emission region in  the HST image ($\sim 1.5\arcsec$ in diameter) which at a distance of 650 pc corresponds 
to 490 AU in radius. We can estimate that the more internal emission zones (for example [O~III]) are located at about 
170~AU from the exciting star. If the assumptions in this work are valid, the extended emitting zone located in the 
direction of the equator is composed of expanding material. We propose that the emission lines are produced in an 
inner extension of the molecular torus located very close to the star, and extending to large enough distances to pinch 
the waist of the lobes. In our scheme the optical emission lines are produced in the inner regions of the torus that 
is surrounding the central star. The high ionization species are located in the innermost region where the gas is being 
shocked by the fast wind, and the less ionized species in the external part. The expansion rate of the neutral species 
is consistent with the low velocity of the CO torus. This scenario corresponds to a very extended torus where the inner 
region is producing the optical emission lines and the radio continuum emission (Kwok et al. 1985), followed by a warm 
H$_2$ component (Smith et al. 2005), that merges into a CO torus (Zweigle et al. 1997), altogether forming a warm dusty 
torus at 260 K (Smith \& Gehrz 2005). In principle, this scenario is in agreement to the wind interaction models 
(Mellema 1993; Franco, Tenorio-Tagle, \& Bodenheimer 1990) at very close distance to the central star, and is compatible with the radio continuum 
spectral index that agrees with mass loss from the central star. The size of the circumstellar structure is small 
and the characteristic times are relatively short (for a structure 490 AU radius the crossing time from the central 
star outwards at a velocity of 30~km~s$^{-1}$ is 80 years, and at the same velocity the distance from the inner border 
to the outer edge of the ionized region can be reached in 53 years). Thus it would be of interest to monitor the kinematic 
behavior of the emission lines, for any possible changes in timescales of decades. Furthermore, we suggest that the group 
of pinched waist planetary nebulae should be investigated to find out if this behavior extends to all of them.

For its part, the core of M1-91 exhibits double peaked Balmer lines profiles. It is proposed that these two components 
can be explained by large optical depth effects in the hydrogen Balmer lines. From the H$\alpha$/H$\beta$ ratio, an 
optical depth of $\tau \sim 4$ was estimated for the H$\alpha$ line. For {M~1-91} we again find a correlation of the 
kinematics of the forbidden lines similar to that of {M~2-9}. We further find a tendency for the atoms with higher ionization 
potential to correspond to denser zones, and to have higher absolute velocities relative to the nucleus, while those 
ions of lower ionization potential are less dense and have lower velocities relative to the nucleus. Although there 
are no observations of a molecular component in {M~1-91}, we extend our interpretation of {M~2-9} to this object by proposing 
that the ionized region is the inner part of an extended torus that has produced such a bipolar morphology. It is estimated 
that the radius of the [S~II] circumstellar emitting region is $\sim 760$~AU (assuming a distance to the nebula of 
1.9 kpc). On the one hand, by assuming a quasi equilibrium mass flow in this region it is found that the distance from 
the nucleus to the inner edge of the circumstellar torus (corresponding to the [Ne III],  [O~III] emitting region) is $\sim 50$~AU.

In conclusion, as has been noted by many authors, M~2-9 and M~1-91 indeed share many characteristics: very young emission 
nebulae of extremely collimated morphology, undisclosed characteristics of the exciting star(s), very dusty circumstellar 
regions, very high density cores that give rise to H$\alpha$ Raman scattering, complex H$\alpha$ profiles, similar 
nuclear spectra with permitted and emission lines that show the same kinematic behavior, and expanding circumstellar tori 
that presumably extend to produce the pinched waist. The sizes of the circumstellar structures are small and the characteristic 
times are relatively short, such  that it would be of interest to monitor the kinematic behavior of the emission lines 
for possible changes that could presumably take place in timescales of decades. Furthermore, we suggest that the cores of 
the group of pinched waist planetary nebulae should be studied spectroscopically to find out if the emission line kinematics 
extends to all of them, and if that is the case, to monitor their behavior in time.


\acknowledgements

ST-P and AA have received support from DGAPA-IN112708, Conacyt-60967 and Conacyt-46904 grants. We are grateful for the 
comments from the referee and many fruitful discussions with L. N. Georgiev.

\vspace{0.2cm}

{\it Note added in proof}: R. Corradi (private communication) kindly informed us that the distance of 650~pc assigned to M~2-9 
(Schwarz et al. 1997) is not correct, and that it should be 1.3~kpc. This value, if confirmed, would modify our size and time 
estimates by a factor of 2.


\begin{thebibliography}{}

\bibitem{} Allen, D.~A., \& Swings, J.~P. 1972, \apj, 174, 583

\bibitem{} Arrieta, A., \& Torres-Peimbert, S. 2000, in ASP Conf. Ser. 199, Asymmetrical Planetary Nebulae II: From Origins to Microstructures, ed. J. H. Kastner, N. Soker, \& S. Rappaport (San Francisco: ASP), 255

\bibitem{} \ibidrule 2003,  ApJS, 147, 97

\bibitem{} Aspin, C., McLean, I.~S., \& Smith, M.~G. 1988, A\&A, 196, 227

\bibitem{} Balick, B. 1989, \aj, 97, 476

\bibitem{} \ibidrule 1999, in ASP Conf. Ser. 188, Optical and Infrared Spectroscopy of Circumstellar Matter, ed. E. Guenther, B. Stecklum, \& S. Klose (San Francisco: ASP), 241

\bibitem{} Balick, B., Icke V., \& Mellema, G., 1997, STScI Press Release PR97-38

\bibitem{} Bachiller, R., Mart\'in-Pintado, J.,  \& Bujarrabal, V.  1990, A\&A, 227, 188 

\bibitem{} Bautista, M.~A. 1999, A\&AS, 137, 529

\bibitem{} \ibidrule 2001, A\&A, 365, 268

\bibitem{} \ibidrule 2004, A\&A, 420, 763

\bibitem{} Bautista, M.~A., Kallman, T.~R., Angelini, L., Liedahl, D.~A., \& Smits, D. P. 1998, \apj, 509, 848

\bibitem{} Bautista, M.~A., Pogge, R.~W., \& DePoy, D.~L. 1995, \apj, 452, 685

\bibitem{} Bautista, M.~A., \& Pradhan, A.~K. 1998,  RevMexAA (SC), 7, 163

\bibitem{} Cahn, J.~H., \& Kaler, J.~B. 1971, \apjs, 22, 319 

\bibitem{} Calvet, N., \& Cohen, M. 1978, \mnras, 182, 687

\bibitem{} Cardelli, J.~A., Clayton, G.~C., \& Mathis, J.~S. 1989, \apj, 345, 245

\bibitem{} Carsenty, U., \& Solf, J.  1983, IAU Symp. 103, Planetary Nebulae, ed. D. R. Flower (Dordrecht: Reidel), 510 

\bibitem{} Corradi, R.~L.~M., \& Schwarz, H.~E. 1995, A\&A, 293, 871

\bibitem{} Doyle, S., Balick, B., Corradi, R.~L.~M., \& Schwarz, H.~E. 2000, \aj, 119, 1339

\bibitem{} Franco, J., Tenorio-Tagle, G., \& Bodenheimer, P. 1990, \apj, 349, 126

\bibitem{} Georgiev, L.~N., Richer, M.~G., Arrieta, A., \& Zhekov, S.~A. 2006, \apj, 639, 185

\bibitem{} Goodrich, R.~W. 1991, \apj, 366, 163

\bibitem{} Grevesse, N., \& Sauval, A.~J. 1998, Space Sci. Rev., 85, 161

\adjustfinalcols

\renewcommand\baselinestretch{1.008}\small

\bibitem{} Guerrero, M.~A., Villaver, E., Machado, A., Garc\'ia-Lario, P., \& Prada, F. 2000, \apjs, 127, 125

\bibitem{} Hora, J.~L., \& Latter, W.~B. 1994, \apj, 437, 281 

\bibitem{} Josselin, E., Bachiller, R., Manchado, A., \& Guerrero, M.~A. 2000, \aap, 353, 363

\bibitem{} Kim, S., Martin, P.~G., \& Hendry, P.~D. 1994, \apj, 422, 164

\bibitem{} Kohoutek, L., \& Surdej, J. 1980, \aap, 85, 161

\bibitem{} Kwok, S.,  Purton, C.~R.,  Matthews, H.~E, \&  Spoelstra, T.~A.~T. 1985, A\&A, 144, 321

\bibitem{} Lee, T.-H., Lim, J., \& Kwok, S. 2007, ApJ, 665, 341

\bibitem{} Livio, M., \& Soker, N. 2001, \apj, 552, 685

\bibitem{} Maciel, W.~J. 1984, A\&AS, 55, 253

\bibitem{} Manchado, A., Guerrero, M.~A., Stanghellini, L., \& Serra-Ricart, M. 1996, The IAC Morphological Catalog of Northern Galactic 
Planetary Nebulae (La Laguna, Espa\~na: Instituto de Astrof\'isica de Canarias) 

\bibitem{} Mathis, J.~S., \& Wallenhorst, S.~G. 1981, \apj, 244, 483

\bibitem{} Mel\'endez, M., \& Bautista, M.~A. 2005, A\&A, 436, 1123

\bibitem{} Mellema, G. 1993, Ph.D. Thesis, University of Leiden

\bibitem{} Osterbrock, D.~E. 1989,  Astrophysics of Gaseous Nebulae and Active Galactic Nuclei (Mill Valley: Univ. Science Books)

\bibitem{} Rodr\'iguez, M., Corradi, R.~L.~M., \& Mampaso, A. 2001, \aap, 377, 1042

\bibitem{} Schmidt, G.~D., \& Cohen, M., 1981, ApJ, 246, 444

\bibitem{} Schwarz, H.~E., Aspin, C., Corradi, R.~L.~M., \& Reipurth, B. 1997, \aap, 319, 267

\bibitem{} Sharpee, B., Williams, R., Baldwin, J.~A., \& van Hoof, P.~A.~M. 2003, \apjs, 149, 157

\bibitem{} Smith, N., Balick, B., \& Gehrz, R.~D. 2005, \aj, 130, 853

\bibitem{} Smith, N., \& Gehrz, R.~D. 2005, \aj, 129, 969

\bibitem{} Solf, J. 1984, \aap, 139, 296

\bibitem{} \ibidrule 2000, \aap, 354, 674

\bibitem{} Swings, J.~P., \& Andrillat, Y. 1979, \aap, 74, 85

\bibitem{} Trammell, S.~R., Goodrich, R.~W., \& Dinerstein, H.~L. 1995, ApJ, 453, 761

\bibitem{} Trammell, S.~R., Kartje, J.~F., Hill, J.~G., \& Dinerstein, H.~L. 1996, in ASP Conf. Ser.  97, Polarimetry of the Interstellar Medium, ed. W. G. Roberge \&  C. B. Doug (San Francisco: ASP), 217

\bibitem{} Walsh, J.~R. 1981, MNRAS, 194, 903

\bibitem{} Zweigle, J., Neri, R., Bachiller, R., Bujarrabal, V., \& Grewing, M. 1997, A\&A, 324, 624


\end{thebibliography}


%\input{torrestab2.tex}

\end{document}

