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Issue
A&A
Volume 545, September 2012
Article Number A75
Number of page(s) 17
Section Extragalactic astronomy
DOI https://doi.org/10.1051/0004-6361/201219689
Published online 10 September 2012

© ESO, 2012

1. Introduction

The stellar and gas distributions in a cluster galaxy can be drastically modified through environmental effects. In Virgo, a young, close (d ~ 17 Mpc) and still dynamically active cluster (see Boselli & Gavazzi 2006, and references therein), the environment effects are clearly seen for the atomic gas. Through tidal interaction and/or ram pressure stripping a spiral galaxy can lose a considerable fraction of its neutral gas and become deficient in HI (Davies & Lewis 1973; Giovanardi et al. 1983; Haynes et al. 1984; Cayatte et al. 1990; Chung et al. 2009a).

For the molecular gas-phase the situation is more complex. The typical density of a molecular cloud core is 105 cm-3 (Caselli et al. 2002), and the molecular gas is more tightly bound in the galaxy potential, being confined to the inner part of the disk which may prevent gas removal via stripping. According to Kenney & Young (1986, 1989), Virgo spirals show no clear signs of deficiency in the molecular component. In the Coma cluster, where the ram pressure is expected to be stronger due to the high ICM densities, there is no clear evidence of molecular deficiency (Casoli et al. 1991). However, as noted by Boselli & Gavazzi (2006), the sample of Casoli et al. (1991) is selected according to IR luminosity, and their results are biased toward strong CO emitters (see also Boselli et al. 1995a, 1997). Using the same data of Kenney & Young (1989), Rengarajan & Iyengar (1992) found evidence of a CO deficit in some Virgo spiral, and Vollmer et al. (2008) detected in NGC 4522 molecular gas displaced from the galactic disk because of ram pressure.

Boselli et al. (2002), Fumagalli & Gavazzi (2008) and Fumagalli et al. (2009) found a molecular depletion in highly HI deficient galaxies, where the HI is depleted well inside the optical disk. Since not all the HI-deficient galaxies examined in their papers are cluster galaxies, it is of interest to analyze a sample of normal and HI-deficient galaxies in a cluster. Also, it is still not clear if the depletion of the molecular component in HI-deficient galaxies is due to ram pressure stripping of the molecular component, or to the lack of HI from which molecular clouds form.

The Herschel Virgo Cluster Survey (HeViCS; Davies et al. 2010; Auld et al. 2012) has mapped an area of 84 square deg2 in the Virgo cluster, and one of the goals of the project is to investigate mechanisms of gas recycling in clusters of galaxies. Cortese et al. (2012) analyzed the galaxies in the Herschel Reference Survey (HRS; Boselli et al. 2010), including many Virgo galaxies observed within HeViCS, and found that the global dust-to-stellar mass ratio decreases as a function of HI deficiency, providing strong evidence of dust stripping. Furthermore they found that at fixed stellar mass, the dust-to-HI mass ratio decreases with the HI deficiency, with a correlation significantly weaker than for the dust-to-stellar mass ratio. They argue that this mild correlation is a consequence of the greater extension of the HI disks which produces a rapid decline of the dust-to-HI ratio in the outer regions of galaxies. For a sample of 35 HeViCS Virgo spirals with available integrated CO fluxes, Corbelli et al. (2012) confirmed that the dust-to-stellar mass ratio decreases and the dust-to-total gas ratio increases when the galaxies are more HI-deficient.

The investigation of the internal properties of galaxies can help understand environmental effects. Cortese et al. (2010) investigated 15 Virgo spiral galaxies, and found that the radial distributions of dust in HI-deficient galaxies are truncated, already in the inner regions of the disks.

It is therefore important to study the relation between the radial distribution of dust and the different gas components within a spiral disk. To reach this goal, it is necessary to have a sample of galaxies observed in all the interstellar medium (ISM) components. The proximity of Virgo allows a good resolution of the gas and dust components (a typical beam size of 20′′ corresponds to a resolution of 1.6 kpc). Several Virgo spiral galaxies have been mapped with the Very Large Array (VLA) in HI, within the survey “VLA Imaging of Virgo spirals in Atomic gas” (VIVA; Chung et al. 2009a). About half of them are mapped in the FIR within the HeViCS project (Davies et al. 2010). However, only a subset of these have observations that trace the spatial distribution of the molecular component, major axis profiles or maps such as those shown in Chung et al. (2007), Kuno et al. (2007). To complete the sample we thus began a project to map the CO emission in the remaining low luminosity objects.

In this paper we present 12CO(1–0) and 12CO(2–1) emission-line observations for seven galaxies, obtained at the IRAM-30 m telescope. We use these new data to investigate the balance between different gas phases in spiral disks and their relation with the dust component as a function of the environment. The paper is organized as follows: in Sect. 2, we describe the sample, and the ancillary data used in the paper. In Sect. 3 we present the new IRAM-30 m observations and the relative results. In Sect. 4 we discuss in detail the radial distribution of the dust and gas phases for the galaxies mapped in our new campaign and for additional Virgo spirals with available CO, HI, and dust maps. We also investigate the variations of the dust-to-gas ratio across the disk. In Sect. 5 we perform a pixel-by-pixel analysis of the relation between the far-infrared (FIR) emission and the gas surface mass density. The main conclusions are summarized in Sect. 6.

2. The sample and ancillary data

We have compiled a sample of bright (mB ≤  13) late-type galaxies in the Virgo cluster with available VIVA and HeViCS observations. Here we present 12CO(1–0) and 12CO(2–1) emission-line maps of four spirals (NGC 4189, NGC 4298, NGC 4299, NGC 4388) obtained with the IRAM-30 m telescope. We also observed a few selected regions of 3 additional low luminosity galaxies, NGC 4351, NGC 4294, and NGC 4424. To enlarge the sample of fully mapped galaxies we also consider eight additional galaxies mapped in 12CO(1–0) at the Nobeyama 45 m telescope by Kuno et al. (2007) at comparable spatial resolution. This sample of 12 mapped galaxies spans a wide range of morphological types, from Sa (NGC 4424) to Scd (NGC 4299). Since we are interested in studying the environmental effects on the gas distribution, we excluded from the sample gravitationally interacting pairs.

Interferometric observations of HI for the VIVA sample are described in Chung et al. (2009a). The velocity resolution of VIVA (10 km s-1) is comparable to the final resolution of our CO data (10.4 km s-1). The sensitivity of the VIVA observations is around 3–5 ×  1019 HI atoms cm-2 (3σ, 0.3 M⊙ pc-2) at a typical angular resolution of ~ 25′′. We quantify the HI deficiency, defHI, in our sample using the prescription of Chung et al. (2009a), which assumes a mean HI mass-diameter relation, independently of morphological type. Our galaxies span a wide range of values, −0.43 ≤  defHI ≤  1.47, revealing different stages of interaction with the environment. Even if we adopt an HI deficiency criterion based on morphological type (e.g. Giovanardi et al. 1983), our sample still covers a wide range of defHI: from −0.2 to 1.1. Given that the uncertainties in defHI are ~0.25–0.3 dex we shall consider a galaxy as non-deficient only if it has defHI ≲ 0.4.

FIR images for dust emission at 100, 160, 250, 350, and 500 μm come from the HeViCS full depth data set (see Auld et al. 2012,for details on the data reduction). The beam full width half maximum for these wavelengths are 9′′, 13′′, 182, 245, and 36′′, respectively.

The dust surface mass density is derived as in Magrini et al. (2011) (see also Smith et al. 2010; Bendo et al. 2010a). The five images in the Herschel photometric bands were convolved to the poorer resolution of the 500 μm, then registered on the same pixel grid. Each pixel was fitted with a single temperature modified black body, with dust emissivity κ ∝ νβ and β = 2; the filter response function was taken into account for the color correction. As shown by Davies et al. (2012) for global fluxes, and by Magrini et al. (2011) for pixel-by-pixel surface brightness analysis, our choice of the dust emissivity provides a satisfactory fit to the spectral energy distribution. If we exclude the notorious uncertainty of the dust emissivity, the uncertainties due to the fitting procedure amounts to 10–15% for the full depth HeViCS data set. Other sources of uncertainty in the determination of the total mass are discussed in Magrini et al. (2011).

In Table 1 we summarize the main properties of the galaxies in the sample, and in Appendix A we describe more in detail the galaxies for which we present new CO observations.

Table 1

General properties of the sample.

3. IRAM 30-m observations

The CO observations were carried out at the IRAM 30-m millimeter-wave telescope on Pico Veleta (Spain) during two observing runs in June 2010 and April 2011.

3.1. Observational setup and data reduction

We observed the 12CO(1–0) and 12CO(2–1) rotational transitions, with rest frequencies of 115.271 and 230.538 GHz, respectively. As front-end we used the Eight MIxer Receiver (EMIR; Carter et al. 2012), and adopted three different back-ends:

  • the 4 MHz filterbank, used for the12CO(1–0) line only, providing a total band widthof 4 GHz, and a channel widthof 4 MHz, i.e.10.4 km s-1;

  • the Wideband Line Multiple Autocorrelator (WILMA) with a bandwidth of 1860 channels (4 GHz), and a spectral resolution of 2 MHz. This corresponds to a spectral resolution of 5.2, and 2.25 km s-1 for the 12CO(1–0) and 12CO(2–1) lines;

  • the VErsatile SPectrometer Assembly (VESPA) for which we selected the resolution of 1.25 MHz, with a smaller bandwidth, 420 MHz. This setup gives a spectral resolution of 3.25, and 1.62 km s-1 for the 12CO(1–0) and 12CO(2–1) lines.

We performed a flux calibration every 15 min, checked the pointing every 2 h, and the focus every 6 h, as customary. The spectra were calibrated in main-beam temperature, assuming the main-beam and forward efficiencies ηmb = 0.78 and ηfor = 0.95 for the 12CO(1–0) and ηmb = 0.59, ηfor = 0.91 for 12CO(2–1), according to the values given in the IRAM-30 m website as of July 2011. The main-beam temperature is , where  is the antenna temperature (Wilson et al. 2009).

To increase the signal-to-noise ratio (S/N) we averaged the spectra taken with all the backends to a common resolution of 10.4 km s-1. We removed the baseline outside the signal window by fitting polynomial functions1 of orders between 1 and 3. The half power beam widths (HPBW) are 225 for 12CO(1–0) and 1125 for 12CO(2–1). We observed our sample in 2 different observing modes: on the fly (OTF) map and position switching (PS). The choice of the observing mode was based on the estimated total CO brightness from FIR emission.

Table 2

Integrated quantities for observations in OTF mode.

Table 3

Integrated quantities for observations in PS mode.

In OTF mode the telescope moves smoothly across the target and collects spectra while scanning. We set the integration time per position to 2 s and the distance between each position along the scan to 4′′, below the critical Nyquist frequency. After each scan a reference spectrum is taken in a region with no sources. Several scans along a chosen direction were taken to cover an area of about 1 square arcmin. The observations were repeated several time along parallel and perpendicular scan directions to reach the required sensitivity. In total, we observed NGC 4189 for 8.1 h, NGC 4298 for 4 h, NGC 4388 for 2.4 h, and NGC 4299 for 4 h. The sensitivities we achieved for the 12CO(1–0) and 12CO(2–1) lines are shown in Cols. 3, 4 of Table 2. The size of each map is given in Col. 2 of Table 2. Scans were done along the RA and Dec directions, with the exception of NGC 4298, where a scan direction parallel and perpendicular to the major axis was used.

The map pixel size is chosen in order to achieve roughly Nyquist sampling, 93 for 12CO(1–0) and 47 for 12CO(2–1) (HPBW/2.4). Spectra at each pixel are summed using the sigma-weighted average to obtain the map. The observations have been smoothed with a Gaussian of size 1/3 of the telescope beam to gain in sensitivity at the expense of a slightly larger beam size (see Mangum et al. 2007). In the following analysis, the 12CO(2–1) maps have been convolved with an appropriate bidimensional Gaussian kernel to match the resolution of the 12CO(1–0).

To estimate the Moment 0 and Moment 1 maps we defined in each spatial pixel a spectral signal window and integrated within it. We calculated the noise σ outside the window and then estimated the noise related to the signal as , where Δv is the width of the window in km s-1 and Nch is the number of channels inside the signal window. In the final cube we took into account only the pixels with S/N ≥ 3.

Three galaxies with expected fainter CO emission were observed in PS mode at a few selected positions, with the telescope switching between the galaxy and an offset position for background subtraction. In these observations we integrated for 15 min on each position and achieved the sensitivities shown in Cols. 3 and 4 in Table 3. The original plan was to observe several positions along the major axis spaced by the 12CO(1–0) beamsize. However, due to bad weather conditions and time limitations this was possible only for NGC 4424, which was observed in 5 positions. For NGC 4294 it was possible to observe only the galaxy center and another position, and finally for NGC 4351 we only observed the galaxy center. The observed positions with respect to the center are shown in Col. 2 of Table 3.

3.2. Results on individual galaxies

Table 2 shows the average value of the integrated intensity for 12CO(1–0) and 12CO(2–1) lines (Cols. 5 and 6), together with the average intensity line ratio (Col. 7) for galaxies observed in OTF mode. In Cols. 8 and 9 we report the average surface brightness and the total mass of molecular hydrogen, obtained integrating over the pixels with surface mass densities above 3σ. We converted the 12CO(1–0) integrated intensities into molecular gas masses using the standard Galactic conversion factor XCO = NH2/ICO = 2 × 1020 cm-2 [K km s-1]-1, estimated in the solar neighborhood (Strong & Mattox 1996; Dame et al. 2001). In Sect. 4 we investigate how the choice of different XCO factors affects the estimate of molecular hydrogen surface mass densities. We estimate the molecular fraction, defined as fmol = MH2/(MH2 + MHI), using the mass of atomic hydrogen given by Chung et al. (2009a).

In Figs. 1, 3, 4 the CO maps obtained in OTF mode are presented in the following order: in the top left and right panels we show the moment 0 and moment 1 map of the 12CO(1–0) intensity emission; in the middle left and right panels the moment 0 12CO(2–1) map and the 12CO(2–1)/12CO(1–0) line intensity ratio are shown. The channel velocity maps with Δv = 10.4 km s-1 are given in the bottom panels.

thumbnail Fig. 1

NGC 4189 maps. Top left: 12CO(1–0) Moment 0 map, contours: 4.2, 6.9, 9.6, 12.3, 15, 17.7 M⊙ pc-2. Top right: 12CO(1–0) Moment 1 with 2000 ≤ v ≤ 2230 and Δv = 28 km s-1. Middle left: 12CO(2–1) Moment 0 map, contours: 4, 7.25, 10.5, 13.75, 17, 20.25, 23.5 M⊙ pc-2. Middle right: 12CO(2–1)/12CO(1–0) ratio. The beam size is illustrated in the bottom left of each panel. Bottom panel: channels map of 12CO(1–0) with 1980 ≤ v ≤ 2250 km s-1 and Δv = 10 km s-1.

thumbnail Fig. 2

Hα emission (from GOLDMine) of NGC 4189 (top row) and NGC 4298 (bottom row) with overplotted 12CO(1–0) (left column), 250 μm emission (middle column), and HI surface mass density (right column). The contours of 12CO(1–0) line are 4.2, 6.9, 9.6, 12.3, 15, and 17.7 M⊙ pc-2 for NGC 4189 and 0.5, 8.8, 17.1, 25.4, 33.7, 42, and 50.3 M⊙ pc-2 for NGC 4298. The 250 μm contours are at 0.23, 0.28, 0.31, 0.36, 0.42, 0.48, 0.56, and 0.65 Jy beam-1 for NGC 4189 and 0.2, 0.3, 0.41, 0.63, 0.84, and 1.06 Jy beam-1 for NGC 4298. The HI surface mass density contours are at 4.8, 7.2, 10, 12.4, 14.8, and 17.2 M⊙ pc-2 for NGC 4189, and at 4.3, 5.2, 6, 7, 8, 9.1, 10.6 M⊙ pc-2 for NGC 4298. The beam size of each observations is shown in the bottom left of each panel.

thumbnail Fig. 3

NGC 4298 maps. Top left: 12CO(1–0) Moment 0 map, contours: 0.5, 8.8, 17.1, 25.4, 33.7, 42, and 50.3 M⊙ pc-2. Top right: 12CO(1–0) Moment 1 with 1035 ≤ v ≤ 1235 and Δv = 20 km s-1. Middle left: 12CO(2–1) Moment 0 map, contours: 0.5, 7.6, 14.8, 21.9, 29, and 36.2 M⊙ pc-2. Middle right: 12CO(2–1) over 12CO(1–0) ratio. The beam size is shown in the bottom left of each panel. Bottom panel: channels map of 12CO(1–0) with 985 ≤ v ≤ 1275 km s-1 and Δv = 10 km s-1.

NGC 4189

The CO maps of NGC 4189 are shown in Fig. 1. The 3σ noise in 12CO(1–0) and 12CO(2–1) maps is 4.2 and 4 M⊙ pc-2, respectively. The total mass of molecular gas is MH2 = 17.4 ± 1.3 × 108 M⊙, and the molecular-to-total gas mass fraction is fmol = 0.43. The 12CO(1–0) emission is less extended (~16 at 4.2 M⊙ pc-2) than the atomic hydrogen (~26 at 4.8 M⊙ pc-2) but comparable to the extent of the optical disk (226). The total intensity of the 12CO(2–1) lines is comparable to that of the 12CO(1–0) line.

The 12CO(1–0) emission (top left panel of Fig. 1) has a regular structure centered in the South-East region and is more homogeneously distributed with respect to the 12CO(2–1) emission, which has a clumpier structure dominated by a few unresolved regions (of 1 pixel size). Taking into account only regions with S/N ≥ 3 the 12CO(2–1)/12CO(1–0) ratio varies between 0.4 and 1 (bottom right panel of Fig. 1). These values are consistent with the bulk of the emission being optically thick.

In the top row of Fig. 2 we compare qualitatively the Hα emission of NGC 4189 (from GOLDMine; Gavazzi et al. 2003) with the 12CO(1–0) (left), HI (middle), and the 250 μm (right) emission. Dust emission at 250 μm is comparable in size with the 12CO(1–0) emission. The distribution of dust emission across the disk is regular and more similar to the CO brightness distribution than to the HI. There is a higher concentration of the dust in the South-East regions, in agreement with a similar feature seen in the HI, 12CO(1–0), and Hα maps.

NGC 4298

The CO maps for NGC 4298 are shown in Fig. 3. We estimate a total mass MH2 = 10.4 ± 0.8 × 108 M⊙ and fmol = 0.65. Both the 12CO(1–0) (top left panel of Fig. 3) and 12CO(2–1) emissions (middle left panel of Fig. 3) peak in the central regions, consistently with the observed Hα emission (bottom row of Fig. 2, from GOLDMine).

The Moment 1 map (top right panel of Fig. 3) shows a quite regular structure in the northern part of the disk and slightly disturbed velocity contours in the South. The 12CO(2–1)/12CO(1–0) ratio has values between 0.3 and 1.4, with the highest values found in isolated clumps of 1–2 pixels in size. The location of these pixels corresponds to the brightest HII regions in the galaxy.

Both the HI and the CO emission have a higher surface mass density in the South-East region, but the 250 μm emission is only slightly asymmetric (see bottom row panels of Fig. 2).

NGC 4388

thumbnail Fig. 4

NGC 4388 maps. Top left panel: 12CO(1–0) Moment 0 map, contours: 1, 12, 23, 34, 45, 56, and 67 M⊙ pc-2. Top right panel: 12CO(1–0) Moment 1 with 2370 ≤ v ≤ 2800 and Δv = 43 km s-1. Middle left panel: 12CO(2–1) Moment 0 map, contours: 0.01, 8.6, 17.3, 26, 34.7, 43.3, and 52 M⊙ pc-2. Middle right panel: 12CO(2–1) over 12CO(1–0) ratio. The beam size is shown in the bottom left of each panel. Bottom panel: channels map of 12CO(1–0) with 2304 ≤ v ≤ 2744 km s-1 and Δv = 10 km s-1.

thumbnail Fig. 5

Top panel: 12CO(1–0) emission of NGC 4388 with overplotted HI emission from VIVA (Chung et al. 2009a). Contours: 15.7, 18.7, 22.1, 25.2, 28.6 M⊙ pc-2. Middle panel: SDSS g-band image of NGC 4388 with overplotted 12CO(1–0) emission, contours: 1, 12, 23, 34, 45, 56, and 67 M⊙ pc-2. Bottom panel: SDSS g-band image of NGC 4388 with overplotted 250 μm dust emission, contours: 0.35, 0.71, 1.07, 1.43, 1.79, and 2.15 Jy beam-1. The beam size is shown in each panel. The arrow in the bottom panel indicates the direction of ram pressure.

Table 4

Dust and gas properties in the sample.

For NGC 4388 the CO maps are shown in Fig. 4. The estimated total mass of molecular hydrogen is 7.1 ± 0.6 × 108 M⊙ and the molecular fraction is fmol = 0.63. Unfortunately the 12CO(2–1) emission has SNR > 3 only in the central regions and the relative map is smaller than 12CO(1–0) map. Both the HI and CO distributions throughout the disk drop sharply below the sensitivity of our observations; this can be seen in the middle panel of Fig. 5, where we overlay the contours of the molecular gas onto a reference g-band image from SDSS-DR6 (Adelman-McCarthy et al. 2008). The atomic and molecular components are truncated at R/R25 ≃ 0.6, where R25 is the radius at a B-band surface brightness of 25 mag arcsec-2 (R25 = a/2, see Col. 8 of Table 1), as expected for highly HI-deficient galaxies (defHI = 1.16). The top panel of Fig. 5 shows that the radial extent of the atomic and molecular phases for this galaxy are comparable.

The structure of the 12CO(1–0) emission is asymmetric, with a more prominent elongated diffuse tail in the east side of the disk (middle panel of Fig. 5). A similar feature is present also in the HI distribution (Chung et al. 2009a), and this overdensity is higher where the gas plume observed in Oosterloo & van Gorkom (2005) connects with the disk. At 250 μm the disk is also slightly asymmetric, similar in shape to the CO distribution (middle and bottom panels of Fig. 5). This asymmetry in the ISM component is probably due to a recent ram pressure episode (see arrow in Fig. 5, Vollmer & Huchtmeier 2003; Pappalardo et al. 2010). Because of the spatial limitation of our map, we were unable to investigate whether there is any stripped molecular gas associated with the elongated HI tail seen by Oosterloo & van Gorkom (2005). The channel maps (top right panel of Fig. 4) do not show the typical “spider diagram” of spiral galaxies, because NGC 4388 is highly inclined (i = 75°) and the isovelocity contours are crowded along the line of sight.

NGC 4299

Because of the low S/N of the observations, the CO emission in this galaxy has been determined by integrating the spectra in the velocity range 130 ≤ v ≤ 300 km s-1 as determined from the HI emission (Chung et al. 2009a), and inside a circle of 35′′ of radius from the galaxy center (Fig. 6). The size of the circle was chosen to match the Hα emission, taking into account all the star forming regions of the disk. The integration gives an intensity ICO = 1.28 K km s-1, with a noise of σ = 0.28 K km s-1 (i.e. SNR ~ 4.5). The estimated mass inside the circle is MH2 = 1.1 ± 0.2 × 108 M⊙ with an average surface mass density of 4 M⊙ pc-2 and one of the lowest molecular mass fractions fmol = 0.08 (Table 4).

NGC 4294, NGC 4351 and NGC 4424

For the galaxies observed in PS mode Table 3 reports the integrated CO line emission (Cols. 5, 6), the line ratio (Col. 7), and the surface mass density (Col. 8) for each position. The total MH2 mass (Col. 9) was estimated assuming an exponentially decreasing CO-intensity for the gas (Kenney et al. 1996). For both NGC 4294 and NGC 4424 we estimated a scalelength of 0.1 R25. Since NGC 4351 was observed only in the central region, we assumed a scalelength 0.2 R25, as determined in Schruba et al. (2011) for a sample of 33 nearby spiral galaxies.

thumbnail Fig. 6

Letf panel: Hα emission of NGC 4299 with overplotted a circle with R = 35′′ selected to integrate the 12CO(1–0) emission. Right panel: averaged spectra of NGC 4299 inside the circle defined in left panel. The vertical shaded lines show the velocity limits used for the integration.

thumbnail Fig. 7

Top panel: SDSS g-band image of NGC 4351. White circles show the beam size for 12CO(1–0) (solid line) and 12CO(2–1) (dashed line). Bottom panel: 12CO(1–0) (black solid line) and 12CO(2–1) (red dashed line) observations in the position showed in the left panel.

Figures 7–9 show the spectra with the associated positions. For NGC 4351 and NGC 4294 we found the lowest H2 masses of the observed sample, 7.1 × 107 M⊙, and 7.3 × 107 M⊙, and also low molecular fractions, fmol = 0.17, and 0.04, respectively. For NGC 4424 we found a MH2 = 4.2 × 108 M⊙, with a molecular fraction fmol = 0.51. The integrated fluxes of both 12CO(1–0) and 12CO(2–1) have comparable values.

3.3. Comparison with previous molecular mass estimates

Figure 10 shows the comparison of our measurement of the molecular gas mass with previous estimates (Stark et al. 1986; Young et al. 1995; Boselli et al. 1995b; Smith & Madden 1997; Chung et al. 2009b). We have rescaled the literature values to match our choice of XCO and distance. Overall, there is a quite good agreement with previous observations. The only exception is NGC 4388 for which Stark et al. (1986) reported MH2 = 2.2 × 109 M⊙, a value ~30% higher than our estimate. The beam size of their telescope is larger, 100′′, and they observed the disk in 5 different positions, with an offset of 15 in East-West and North-South directions. Surprisingly they found in the East and West side of the disk an emission comparable to the value in the center of the disk. Our map does not extend farther than 80′′ along the disk; however the emission decreases steeply already at ~30–40′′ from the center. Our molecular mass estimate agrees with the value of Kenney & Young (1986) which have sampled similar regions of this galaxy as Stark et al. (1986), but with a smaller beam.

Cortés et al. (2006) carried out interferometric observations of the 12CO(1–0) emission line of the central regions of NGC 4424, and found a molecular gas mass of 1.6 × 108 M⊙. Given the larger area covered by our observations, this is compatible with the mass value we find in the central beam and with our higher total mass value.

Our observations have a higher sensitivity and better spatial resolution than previous ones. For most of the galaxies, σ1−0 is between 4 and 11 mK. Moreover all the galaxies observed in OTF mapping mode have not been mapped previously.

4. Environmental effects on the radial distribution of gas and dust

thumbnail Fig. 8

Same as Fig. 7 for the two positions observed in NGC 4294. The RA and Dec offsets with respect to the center position are indicated in the top right of each spectrum panel.

thumbnail Fig. 9

Same as Fig. 8 for the five positions observed in NGC 4424.

In this section we study the influence of the environment on the surface mass density distributions for the molecular, atomic, and dust components. We consider the new mapped galaxies together with 8 additional galaxies observed at the Nobeyama 45-m telescope (Kuno et al. 2007). The beam size of this set of observations is 15′′, comparable to the IRAM-30 m telescope beam (225 for 12CO(1–0) line), but the typical rms noise level is 40–100 mK, a factor of ten higher than the level achieved by our observations.

We divide our galaxies in three groups: non-deficient galaxies with defHI < 0.4 (NGC 4299, NGC 4254, NGC 4189, NGC 4321), galaxies slightly HI-deficient, with 0.4  ≤  defHI  ≤  0.7 (NGC 4298, NGC 4535, NGC 4192, NGC 4501), and HI-deficient galaxies, with defHI > 0.7 (NGC 4402, NGC 4579, NGC 4388, NGC 4569). We use this classification to characterize the environmental effects in our sample.

We consider the total gas surface density as the sum of the atomic and molecular gas component, since we did not take into account the helium contribution. To scale the quoted surface mass density for helium, one should multiply the recovered values by a factor ~1.36. We also neglected the ionized gas, which is expected to be a small fraction of the total gas mass.

4.1. Radial profiles

The radial profiles are obtained averaging surface mass density in elliptical annuli of width 0.15 R/R25, using the inclination and the major axis position angle given in Cols. 10, 11 of Table 1. The average value in each annulus is plotted when the total number of pixels with surface mass density above the threshold value is at least 20% of the total number of pixels inside the annulus. Finally the radial profiles are corrected to face-on values.

Figure 11 shows the radial profiles of the gas and dust components. The panels are ordered according to the defHI parameter of each galaxy: the top, middle, and bottom rows refer to the non-deficient, slightly HI-deficient, and HI-deficient galaxies, respectively. The open blue circles refer to the molecular surface mass density for a constant Galactic XCO value. The radial profile of the atomic component is plotted using red star symbols. Since the beam size of the HI observations is similar to that of the CO observations, we kept all images at their original spatial resolution (~20–25′′). The total gas profile is shown with black diamonds. The triangles connected by the green solid line show the dust surface mass density at the resolution of 36′′.

As a general trend, the molecular gas has a higher surface mass density than the HI gas in the inner regions. In the external regions the mass surface density of H2 decreases, and the relation between the total gas and the HI distribution depends on the environmental effects undergone by the galaxy. The dominance of the molecular gas in central regions is consistent, for some galaxies in our sample, with the results obtained in a sample of 33 nearby galaxies by Bigiel et al. (2008). By analyzing the surface mass density of the different gas phases they recovered a threshold at which most of the gas is converted into molecular form, i.e. 9 M⊙ pc-2 (see also Krumholz et al. 2008, 2009, for theoretical modeling of this effect). In Fig. 11 the horizontal dotted line marks this value of surface mass density. In the central regions of NGC 4535, NGC 4501, and NGC 4579 the molecular phase is already dominant where the HI surface mass density is only 1 M⊙ pc-2. However, these are HI-deficient galaxies and variations of the threshold value may occur. NGC 4299, the galaxy with the lowest metallicity in our sample, shows a lack of CO. In the central regions the HI dominates the total gas surface mass density, even if for this galaxy it is above the threshold value.

thumbnail Fig. 10

Comparison of our results with previous observations. References: Young et al. (1995) (blue triangles), Chung et al. (2009b) (red star), Stark et al. (1986) (black circles), Smith & Madden (1997) (magenta squares), and Boselli et al. (1995b) (green diamonds). Arrows represents upper or lower limits.

The estimate of the molecular gas profile strongly depends on the assumption for the XCO factor. We have used so far the Galactic solar neighborhood value, but there is ample evidence that XCO varies in other galaxies as a function of physical conditions (e.g. Maloney & Black 1988; Israel 1997; Barone et al. 2000; Boselli et al. 2002; Israel et al. 2003; Israel 2005; Bolatto et al. 2008; Narayanan et al. 2005; Leroy et al. 2011; Shetty et al. 2011; Narayanan et al. 2012; Bell et al. 2007). Thus, we consider an alternative conversion factor, which depends on the metallicity, and varies radially following the metallicity gradient along the disk: (1)with Zc equals to the oxygen abundance in central regions, expressed in the usual form 12+log O/H. Since we do not have Zc and gradients for all the galaxies in our sample we assume the central metallicity estimated using the mass metallicity relation (Tremonti et al. 2004) and reported in Col. 12 of Table 1. We fix the solar value at Z⊙ = 8.69, as recently found by Asplund et al. (2009), and assume a radial metallicity gradient ∝ − 0.4 R/R25, as found by Magrini et al. (2011) for a sample of Virgo spirals. This slope should reproduce the effect of a metallicity gradient.

The blue dashed lines in Fig. 11 represent the surface molecular mass density calculated according to Eq. (1). The choice of a radially varying XCO as a function of metallicity affects significantly the radial profiles and can reduce the molecular mass content up to 60%. Considering for our sample a typical metallicity Zc = 9.09 and according to Eq. (1), we obtain for the central regions XCO = 0.78 × 1020 [K km s-1]-1, considerably lower than the Galactic XCO. For NGC 4189, if we use the galactic XCO we obtain a centrally peaked molecular distribution with a surface mass density 2–3 times larger than HI in the central region. For R/R25 ≳ 0.7, when the surface mass density is below 4 M⊙ pc-2, the HI phase dominates over the molecular phase. Rescaling the conversion factor as in Eq. (1), because of the high value of Zc = 9.09, the atomic hydrogen becomes the dominant component already at R/R25 ≳ 0.3, and the surface mass density of the molecular gas is never above 6 M⊙ pc-2. For NGC 4298 the molecular gas in the central region is always dominant for every choice of the conversion factor, and the effect of a rescaling is to reduce the surface mass density value according to the metallicity. For both NGC 4189 and NGC 4298 the effect of a radially varying conversion factor is negligible in the external regions, R/R25 > 0.6, where, according to our assumed gradient, the metallicity approaches solar. For NGC 4388 the molecular component with a variable XCO has surface mass densities comparable to the atomic component, as a consequence of the strong HI-deficiency (defHI = 1.16). For NGC 4299 (top left panel of Fig. 11) the effect on the total gas is negligible, because the galaxy is HI dominated (defHI = −0.43).

The radial profiles of dust surface mass densities are shown as green triangles in Fig. 11. The slope of the dust profiles tends to be intermediate between the molecular component (steeper slope) and the atomic component (flatter).

To characterize the environmental effects in our sample, we show in Fig. 12 the bin-averaged profiles of the dust, atomic, molecular, and total gas surface mass densities for the three subsamples defined above. In each panel red circles, black triangles, and blue stars, indicate the non-deficient, slightly HI-deficient, and HI-deficient subsamples. For the molecular component we show only the case for the Galactic XCO, since the conclusions about the environmental effects are similar even if the XCO defined in Eq. (1) is used.

thumbnail Fig. 11

Dust (green triangles), atomic gas (red stars), and molecular gas (blue circles) surface mass densities as a function of R/R25 for the extended sample. The molecular component has been estimated using the Galactic XCO (solid line) and the radially varying XCO (Eq. (1); dashed line). The total gas mass surface density (for the Galactic XCO only), is indicated with black diamonds. The horizontal dotted line indicates Σ = 9 M⊙ pc-2.

Galaxies with higher HI deficiency have steeper gas profiles and lower surface mass densities, indicating that the lack of atomic hydrogen affects the formation of the molecular gas, and thus the total gas component. However this trend is not completely obvious because, as shown by Fumagalli et al. (2009), HI deficiency is a condition necessary but not sufficient for molecular gas deficiency. In our subsample of HI-deficient galaxies, according to the definition of H2-deficiency of Fumagalli et al. (2009), only NGC 4579 is strongly H2-deficient, and for this reason the trend observed for the atomic hydrogen in the bottom right panel of Fig. 12 is less evident for the molecular gas component (top left panel in Fig. 12).

thumbnail Fig. 12

Average radial surface mass densities of the dust (bottom left panel), ΣHI (bottom right panel), ΣH2 (top left panel), and ΣHtot (top right panel). Red circles, black triangles, and blue stars represent the average profiles for non-deficient (defHI < 0.4), slightly HI-deficient (0.4 ≤  defHI ≤  0.7), and HI-deficient (defHI > 0.7) galaxies, respectively. The molecular component has been estimated using the Galactic XCO.

As the HI deficiency increases, the surface mass densities are lower for all the ISM components. However, for the atomic hydrogen in the central regions of the disks of HI-deficient galaxies, the surface mass density appears somewhat higher than that of slightly HI-deficient galaxies. This effect could be understood as ram-pressure-induced compression of the gas in the inner region of the disk (Tonnesen & Bryan 2009) that increases locally the gas density (Byrd & Valtonen 1990). Alternatively, the differences in surface mass density can be simply due to uncertainties in the inclination corrections for two of the HI-deficient galaxies, NGC 4388 and NGC 4402, which are nearly edge-on.

Dust radial profiles of HI-deficient galaxies (blue stars) decrease steeply at R/R25 >  0.4, with respect to the non-deficient galaxies (red circles), but have comparable radial slopes in the inner parts of the disk, since the dust is more confined to the disk where ram pressure is less efficient in the ISM removal. Such a trend has already been found by Cortese et al. (2010).

4.2. Dust-to-gas ratios

In Fig. 13 we show the radial profiles of the dust-to-gas ratio for the three subsamples of galaxies, obtained after convolving the gas components to the resolution of the dust mass maps. The black solid line in each panel shows the bin-averaged profiles for each subset.

thumbnail Fig. 13

Dust-to-gas ratio as a function of R/R25 for non-deficient (red circles), defHI < 0.4, slightly HI-deficient 0.4 ≤  defHI ≤  0.7, and HI-deficient galaxies (blue stars), defHI > 0.7. The solid line shows the profile for the new mapped galaxies. The black solid line shows the average profiles. The molecular content has been estimated using both the Galactic XCO (top row) and the XCO of Eq. (1) (bottom row).

In the top panels of Fig. 13 we use the constant Galactic XCO. The averaged profiles show only moderate change with HI deficiency; the dust-to-gas ratio is almost constant throughout the disk for non-deficient galaxies (top left panel of Fig. 13), while it increases with radius for the strongly HI-deficient subset (top right panel of Fig. 13). This could be explained because the gas, particularly the atomic component, is more easily stripped with respect to the dust (Cortese et al. 2012), thus producing a rapid increase of the dust-to-gas ratio at R/R25 > 0.2. However, the average slope of this subset may be biased by the highly HI-deficient galaxy NGC 4569 (defHI = 1.47), which has a gas profile that decreases rapidly at small radii.

Similar trends with HI deficiency are found when the metallicity dependent XCO is used (bottom panels of Fig. 13). However, there is a systematic change in the dust-to-gas ratio. Because of the lower XCO in the metal-rich nuclei, the H2-dominated total gas mass surface densities decrease. Thus, the central dust-to-gas ratio increases, consequently producing negative radial gradients. As shown in the bottom left panel of Fig. 13, this effect is more evident for non-deficient objects, confirming what was already found by Magrini et al. (2011) and by Foyle et al. (2012) on the nearby galaxy M 83. Also, Muñoz-Mateos et al. (2009), using Spitzer data, find decreasing dust-to-gas profiles for an XCO factor that depends on the metallicity gradient. Simple considerations imply that the dust-to-gas ratio should follow linearly the metallicity; both dust-to-gas and metallicity trace metals, the former in the solid and the latter in the gaseous phase (Draine et al. 2007). Thus, a lower-than-Galactic XCO in a galaxy’s center appears to be necessary for the dust-to-gas ratio to have the same gradient observed for the metallicity. However, because of the large uncertainties in the absolute calibration of the metallicity indicators (e.g. Magrini et al. 2011), this statement is inconclusive. Indeed, independently of the metallicity, when the surface mass density is dominated by the atomic gas at all radii the dust-to-gas ratio decreases with radius, similarly to the metallicity gradient (see e.g. Bendo et al. 2010b, for NGC 2403; and Fritz et al. 2012; Smith et al. 2012, for the Andromeda galaxy).

5. The linear relation between Σgas and F250: a pixel by pixel analysis

Figure 14 shows the surface mass density of H2 (top panel, using a Galactic XCO), HI (middle panel), and the total gas (bottom panel) as a function of the surface brightness at 250 μm images (F250) on a pixel-by-pixel basis. We chose this band because of its good sensitivity and better spatial resolution. Moreover, the surface brightness at one wavelength is preferred over the dust surface mass density because it is related to directly observable quantities; the scatter in the relation increases if the dust surface mass density is used instead of the 250 μm surface brightness. Because of the poorer resolution of some of the HI maps, we smoothed all the data to a beam size of 36′′, and regridded all the maps to the same pixel size. Each pixel covers an area of ~9 kpc2, except for NGC 4189 (at a distance of 32 Mpc), for which the area corresponds to ~30 kpc2. The top panel of Fig. 14 shows that ΣH2 is proportional to the surface brightness at 250 μm. Considering all pixels above the 250 μm surface brightness 3σ threshold value of 0.07 Jy beam-1 and ΣH2 ≥ 3σ = 1.15 M⊙ pc-2, we found a best linear fit: (2)with a Pearson correlation coefficient r = 0.86.

The relation is tighter for non-deficient galaxies (red circles) with respect to galaxies with a pronounced HI-deficiency (blue stars). When computing a best linear fit we found a slope m = 1.29 ± 0.07 (r = 0.91) for the non-deficient subsample, and m = 1.08 ± 0.1 (r = 0.81) for the HI-deficient one. For both subsamples we find a similar intercept, ~1.6. The slightly better correlation between the 250 μm surface brightness and ΣH2 for non-deficient objects could indicate that the environment perturbs both the molecular and dust disk, but in somewhat different ways, thus increasing the scatter.

thumbnail Fig. 14

Surface mass densities of molecular hydrogen (top panel), atomic hydrogen (middle panel), and total gas (bottom panel) as a function of the surface brightness at 250 μm on a pixel-by-pixel basis. Different symbols distinguish non-deficient (red circles), slightly HI-deficient (black triangles), and HI-deficient galaxies (blue stars). The pixel size has been set to 36′′. The horizontal dotted lines show the 3σ level for the molecular and total gas (3σ = 1.15 M⊙ pc-2). The diagonal lines show the best linear fits. ΣH2 in the top and bottom panels has been estimated using the constant Galactic XCO. A typical error bar is shown in the bottom panel.

In the middle panel of Fig. 14 no clear correlation can be seen between ΣHI and F250. The spread is larger than for ΣH2 and it appears that non-deficient galaxies have higher ΣHI (~5 M⊙ pc-2 for F250 ~ 0.3 Jy beam-1) than HI-deficient ones (~0.5 M⊙ pc-2), though both samples contain outliers. This indicates that uncertainties in the defHI parameter prevent a clear characterization of the environment effects on all galactic properties.

Adding the HI contribution to all pixels with detected molecular gas, as shown in the bottom panel of Fig. 14, we find: (3)The slope in this case is shallower than that in Eq. (2) and roughly linear. However, the goodness of the correlation does not change significantly when adding HI (r = 0.87). This is in contrast with the correlation found by Corbelli et al. (2012) between the integrated properties: the total gas mass correlates better globally with the dust mass than the single gas components. A similar discrepancy is seen in the relations between star formation rate and gas mass surface densities; while a strong correlation is found between the global star formation rates and the total gas mass surface densities (Kennicutt 1989), a resolved analysis gives a better correlation with the molecular gas surface density (Bigiel et al. 2008; Schruba et al. 2011).

If we use an XCO that varies radially according to Eq. (1), the correlation between Σgas and F250 is slightly degraded (r = 0.8), and the slope becomes shallower (m = 0.75 ± 0.04). The radial variation of the XCO produces in fact flatter radial profiles for the molecular surface mass density, which explains the shallower slope.

Eales et al. (2012) proposed that dust emission can be used to map the ISM in high redshift galaxies: adopting the emissivity of Galactic dust, they derived a linear correlation between the monochromatic luminosity at 250 μm and the total mass of hydrogen (using the Galactic XCO and calibrating the relation on a few nearby galaxies observed with Herschel and with gas mass estimates). With the dust temperatures measured on our maps (T ~ 20 K, see Davies et al. 2012), their correlation can be converted to log  Σgas = log  F250 + 1.62, in good agreement with Eq. (3). However, we had to assume that the 250 μm luminosity arises from the same region occupied by gas, which is guaranteed for the molecular gas component only.

6. Conclusions

In this paper we present new IRAM-30 m 12CO(1–0) and 12CO(2–1) observations of a sample of spiral galaxies in the Virgo cluster which are in the HeViCS field and in the VIVA sample. We observed four galaxies in OTF mode and three in PS mode. For observations done in OTF mode CO-maps are discussed; for galaxies observed in PS mode we only estimate the total molecular content. We include in our analysis a sample of eight galaxies observed in 12CO(1–0) by Kuno et al. (2007), and investigate the effects of the environment on the surface mass density of the atomic and molecular gas phases and on the dust. Maps of the dust surface mass density have been estimated using the HeViCS data set (Davies et al. 2010; Auld et al. 2012). Finally we analyze the relation between the different gas components and the dust distributions in relation to the defHI parameter. Our main conclusions are the following:

  • 1.

    For the new observed galaxies the global molecu-lar gas-to-total mass fraction in the sample varies in therange 0.04 ≤ fmol ≤ 0.65, using the Galactic XCO, with the undisturbed galaxies showing the lowest values. Disturbed galaxies have a lower atomic and molecular content but the environment is more efficient at removing the HI than the molecular gas.

  • 2.

    The total gas surface mass density in the central regions of the disk, above a threshold of 9 M⊙ pc-2 is preferentially found in molecular form, in agreement with the results of Bigiel et al. (2008). However, a few HI-deficient galaxies show an order of magnitude lower threshold, indicating that the value may depend on local conditions. Highly HI-deficient galaxies have H2, HI, and dust surface mass densities in the the central region lower by about a factor 2, suggesting gas and dust removal even in the inner disks.

  • 3.

    Galaxies with high HI deficiency have steeper radial profiles of the total gas, as a consequence of the combined effect of HI and H2 deficiency, in agreement with previous studies (Fumagalli et al. 2009; Cortese et al. 2012). The average radial profiles of the dust surface mass density in HI-deficient galaxies is steeper than in HI-normal galaxies. This is consistent with the results found by Cortese et al. (2010), on the radial distribution of the flux density at 250, 350 and 500 μm.

  • 4.

    The dust-to-gas ratio for non-deficient galaxies has a radial trend which depends on the XCO conversion factor, as already discussed by Magrini et al. (2011). This ratio decreases radially only if XCO increases radially due to metallicity gradients. For HI-deficient galaxies instead, the dust-to-gas ratio stays constant or increases radially, independently of the XCO factors used. This indicates that atomic gas is stripped more efficiently than the dust in a cluster environment. Since both dust and gas radial profiles truncate sharply well inside the optical disk, the dust-to-gas ratio cannot be traced at large galactocentric radius in highly perturbed galaxies.

  • 5.

    We observe a tight pixel-by-pixel correlation between the mass surface density of the molecular hydrogen and the 250 μm surface brightness. In HI-deficient galaxies the correlation weakens, suggesting that the environment is perturbing both gas and dust. Adding the atomic gas to the molecular component makes the correlation with the 250 μm surface brightness linear, but the correlation coefficient does not change significantly. Thus it is not clear if the dust emission at long FIR wavelengths can be used as a tracer of the total gas surface mass density, or of the molecular gas only.

1

Most of data reduction was performed using GILDAS software: http://www.iram.fr/IRAMFR/GILDAS

Acknowledgments

We thank the referee J. Kenney for his useful comments which lead to substantial improvement of the original manuscript. We would like to thank D. Munro for freely distributing his Yorick programming language (available at http://www.maumae.net/yorick/doc/index.html). This work has benefited from research funding from the European Community’s sixth Framework Programme (RadioNet). C.P., L.M., S.B., L.H., E.C., C.G., S.diS.A., are supported through the ASI-INAF agreement I/016/07/0 and I/009/10/0. L.H. and C.P. are also supported by a PRIN-INAF 2009/11 grant. The research leading to these results has received funding from the European Community’s Seventh Framework Programme (/FP7/2007-2013/) under grant agreement No 229517. C.P. would like to thank Fabrizio Massi, and Riccardo Cesaroni for their help during data reduction. C.P. also thanks Christof Buchbender for the really productive “tapas y cerveza” night in Granada.

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Appendix A: Notes on individual objects

A.1. NGC 4189

NGC 4189 is a relatively undisturbed Sc late-type galaxy (defHI = 0.25) located at the edge of the Virgo cluster. The determination of its distance and the membership to the cluster is debated, because Virgo is a cluster that is still forming and comprises distinct clouds of galaxies at different distances. Using the relation of Tully & Fisher (1977) in B-band, Yasuda et al. (1997) estimated for NGC 4189 a distance of 33.8 Mpc, i.e. more than twice the distance of the Milky Way from the center of the cluster (~17 Mpc). Solanes et al. (2002), starting from the HI distribution, reconstructed the structure of the cluster, and found galaxies with elongated gas-tails, due to environmental effects, out to cluster centre distances of 25–30 Mpc. These results imply that at the distance of NGC 4189 from the cluster center it is still possible to find galaxies that are Virgo members. In the following we considered for NGC 4189 the average distance of 32 Mpc for the sub-group to which NGC 4189 belongs, determined using the Tully-Fisher relation in H-band (Gavazzi et al. 1999). The HI disk is more extended than the optical disk, showing a symmetric warp (Chung et al. 2009a) with enhanced emission in the South-East regions, in correspondence to a ridge of Hα emission (Koopmann & Kenney 2004).

A.2. NGC 4298

NGC 4298 is an Sc galaxy (defHI = 0.41) located in the North-East region of Virgo at a projected distance of 3.2° from M 87. The galaxy has flocculent spiral structure (Elmegreen et al. 1999) and has a companion, NGC 4302, at a projected distance of 11 kpc, having a comparable velocity (Δv ≈ 30 km s-1). Both galaxies show HI-tails pointing away from M 87, probably due to ram pressure stripping and/or tidal interaction. However the undisturbed stellar component of NGC 4302 favors the ram pressure hypothesis (see Chung et al. 2007). The HI disk in NGC 4298 is compressed in the South-East region, in the direction of M87, as a consequence of the ram pressure exerted by the intracluster medium. The Hα emission is mostly concentrated in the central region of the disk, except for an enhanced arc in South-East direction, in correspondence with the HI compressed region (Koopmann & Kenney 2004).

A.3. NGC 4299 and NGC 4294

NGC 4299 is an Scd galaxy at a distance of 2.5° from M87 (defHI = −0.43) which forms a possible interacting pair with NGC 4294 (defHI = −0.11). However both galaxies retain their HI gas which forms an extended HI disk (Chung et al. 2009a) with a long HI tail to the South West. Koopmann & Kenney (2004) showed that this pair has enhanced Hα emission in the northern part of the disk, where the HI is compressed. This indicates a stronger star formation rate, probably due to tidal interaction.

A.4. NGC 4351

NGC 4351 is an Sc galaxy which is close to the cluster centre, about 1.7° from M 87 and is slightly HI-deficient (defHI = 0.23). It shows a compression in the HI emission in the north eastern region, due to ram pressure (Chung et al. 2009a). The brightest optical region is offset with respect to the center, and also the Hα emission is asymmetric with enhanced emission in a circumnuclear region (Koopmann & Kenney 2004).

A.5. NGC 4388

NGC 4388 is an Sab galaxy located at a projected distance of 1.3° from M87. The galaxy is one of the most HI-deficient in our sample (defHI = 1.16), showing HI distribution truncated within the optical disk (Chung et al. 2009a). Yoshida et al. (2002) found that the Hα emission extends for 35 kpc in the north east direction. Vollmer & Huchtmeier (2003) found HI gas out to 20 kpc along the same direction as the Hα emission. Further observations showed that the HI tail is even more extended, 110 kpc, with a mass of 3.4 × 108 M⊙ (Oosterloo & van Gorkom 2005). The most probable origin of this elongated tail is a ram pressure stripping episode that occurred about 200 Myr ago (Pappalardo et al. 2010).

A.6. NGC 4424

NGC 4424 is an Sa galaxy with defHI = 0.97. Chung et al. (2007) showed a long HI tail pointing away from M 87 and Kenney et al. (1996) observed star forming regions in the central kpc, but no star formation has been detected beyond these regions. Cortés et al. (2006) found two lobes in the molecular gas component in the central 3 kpc, with total mass MH2 ~ 4.0 × 107 M⊙. The central regions are highly disturbed and asymmetric in the radio continuum, probably as a consequence of strong gravitational interaction (Chung et al. 2009a).

All Tables

Table 1

General properties of the sample.

In the text
Table 2

Integrated quantities for observations in OTF mode.

In the text
Table 3

Integrated quantities for observations in PS mode.

In the text
Table 4

Dust and gas properties in the sample.

In the text

All Figures

thumbnail Fig. 1

NGC 4189 maps. Top left: 12CO(1–0) Moment 0 map, contours: 4.2, 6.9, 9.6, 12.3, 15, 17.7 M⊙ pc-2. Top right: 12CO(1–0) Moment 1 with 2000 ≤ v ≤ 2230 and Δv = 28 km s-1. Middle left: 12CO(2–1) Moment 0 map, contours: 4, 7.25, 10.5, 13.75, 17, 20.25, 23.5 M⊙ pc-2. Middle right: 12CO(2–1)/12CO(1–0) ratio. The beam size is illustrated in the bottom left of each panel. Bottom panel: channels map of 12CO(1–0) with 1980 ≤ v ≤ 2250 km s-1 and Δv = 10 km s-1.
In the text
thumbnail Fig. 2

Hα emission (from GOLDMine) of NGC 4189 (top row) and NGC 4298 (bottom row) with overplotted 12CO(1–0) (left column), 250 μm emission (middle column), and HI surface mass density (right column). The contours of 12CO(1–0) line are 4.2, 6.9, 9.6, 12.3, 15, and 17.7 M⊙ pc-2 for NGC 4189 and 0.5, 8.8, 17.1, 25.4, 33.7, 42, and 50.3 M⊙ pc-2 for NGC 4298. The 250 μm contours are at 0.23, 0.28, 0.31, 0.36, 0.42, 0.48, 0.56, and 0.65 Jy beam-1 for NGC 4189 and 0.2, 0.3, 0.41, 0.63, 0.84, and 1.06 Jy beam-1 for NGC 4298. The HI surface mass density contours are at 4.8, 7.2, 10, 12.4, 14.8, and 17.2 M⊙ pc-2 for NGC 4189, and at 4.3, 5.2, 6, 7, 8, 9.1, 10.6 M⊙ pc-2 for NGC 4298. The beam size of each observations is shown in the bottom left of each panel.
In the text
thumbnail Fig. 3

NGC 4298 maps. Top left: 12CO(1–0) Moment 0 map, contours: 0.5, 8.8, 17.1, 25.4, 33.7, 42, and 50.3 M⊙ pc-2. Top right: 12CO(1–0) Moment 1 with 1035 ≤ v ≤ 1235 and Δv = 20 km s-1. Middle left: 12CO(2–1) Moment 0 map, contours: 0.5, 7.6, 14.8, 21.9, 29, and 36.2 M⊙ pc-2. Middle right: 12CO(2–1) over 12CO(1–0) ratio. The beam size is shown in the bottom left of each panel. Bottom panel: channels map of 12CO(1–0) with 985 ≤ v ≤ 1275 km s-1 and Δv = 10 km s-1.
In the text
thumbnail Fig. 4

NGC 4388 maps. Top left panel: 12CO(1–0) Moment 0 map, contours: 1, 12, 23, 34, 45, 56, and 67 M⊙ pc-2. Top right panel: 12CO(1–0) Moment 1 with 2370 ≤ v ≤ 2800 and Δv = 43 km s-1. Middle left panel: 12CO(2–1) Moment 0 map, contours: 0.01, 8.6, 17.3, 26, 34.7, 43.3, and 52 M⊙ pc-2. Middle right panel: 12CO(2–1) over 12CO(1–0) ratio. The beam size is shown in the bottom left of each panel. Bottom panel: channels map of 12CO(1–0) with 2304 ≤ v ≤ 2744 km s-1 and Δv = 10 km s-1.
In the text
thumbnail Fig. 5

Top panel: 12CO(1–0) emission of NGC 4388 with overplotted HI emission from VIVA (Chung et al. 2009a). Contours: 15.7, 18.7, 22.1, 25.2, 28.6 M⊙ pc-2. Middle panel: SDSS g-band image of NGC 4388 with overplotted 12CO(1–0) emission, contours: 1, 12, 23, 34, 45, 56, and 67 M⊙ pc-2. Bottom panel: SDSS g-band image of NGC 4388 with overplotted 250 μm dust emission, contours: 0.35, 0.71, 1.07, 1.43, 1.79, and 2.15 Jy beam-1. The beam size is shown in each panel. The arrow in the bottom panel indicates the direction of ram pressure.
In the text
thumbnail Fig. 6

Letf panel: Hα emission of NGC 4299 with overplotted a circle with R = 35′′ selected to integrate the 12CO(1–0) emission. Right panel: averaged spectra of NGC 4299 inside the circle defined in left panel. The vertical shaded lines show the velocity limits used for the integration.
In the text
thumbnail Fig. 7

Top panel: SDSS g-band image of NGC 4351. White circles show the beam size for 12CO(1–0) (solid line) and 12CO(2–1) (dashed line). Bottom panel: 12CO(1–0) (black solid line) and 12CO(2–1) (red dashed line) observations in the position showed in the left panel.
In the text
thumbnail Fig. 8

Same as Fig. 7 for the two positions observed in NGC 4294. The RA and Dec offsets with respect to the center position are indicated in the top right of each spectrum panel.
In the text
thumbnail Fig. 9

Same as Fig. 8 for the five positions observed in NGC 4424.
In the text
thumbnail Fig. 10

Comparison of our results with previous observations. References: Young et al. (1995) (blue triangles), Chung et al. (2009b) (red star), Stark et al. (1986) (black circles), Smith & Madden (1997) (magenta squares), and Boselli et al. (1995b) (green diamonds). Arrows represents upper or lower limits.
In the text
thumbnail Fig. 11

Dust (green triangles), atomic gas (red stars), and molecular gas (blue circles) surface mass densities as a function of R/R25 for the extended sample. The molecular component has been estimated using the Galactic XCO (solid line) and the radially varying XCO (Eq. (1); dashed line). The total gas mass surface density (for the Galactic XCO only), is indicated with black diamonds. The horizontal dotted line indicates Σ = 9 M⊙ pc-2.
In the text
thumbnail Fig. 12

Average radial surface mass densities of the dust (bottom left panel), ΣHI (bottom right panel), ΣH2 (top left panel), and ΣHtot (top right panel). Red circles, black triangles, and blue stars represent the average profiles for non-deficient (defHI < 0.4), slightly HI-deficient (0.4 ≤  defHI ≤  0.7), and HI-deficient (defHI > 0.7) galaxies, respectively. The molecular component has been estimated using the Galactic XCO.
In the text
thumbnail Fig. 13

Dust-to-gas ratio as a function of R/R25 for non-deficient (red circles), defHI < 0.4, slightly HI-deficient 0.4 ≤  defHI ≤  0.7, and HI-deficient galaxies (blue stars), defHI > 0.7. The solid line shows the profile for the new mapped galaxies. The black solid line shows the average profiles. The molecular content has been estimated using both the Galactic XCO (top row) and the XCO of Eq. (1) (bottom row).
In the text
thumbnail Fig. 14

Surface mass densities of molecular hydrogen (top panel), atomic hydrogen (middle panel), and total gas (bottom panel) as a function of the surface brightness at 250 μm on a pixel-by-pixel basis. Different symbols distinguish non-deficient (red circles), slightly HI-deficient (black triangles), and HI-deficient galaxies (blue stars). The pixel size has been set to 36′′. The horizontal dotted lines show the 3σ level for the molecular and total gas (3σ = 1.15 M⊙ pc-2). The diagonal lines show the best linear fits. ΣH2 in the top and bottom panels has been estimated using the constant Galactic XCO. A typical error bar is shown in the bottom panel.
In the text

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