Environmental cluster effects and galaxy evolution: The Hi properties of the Abell clusters A85/A496/A2670

In this work we study the Abell clusters A85 ($z=$ 0.055), A496 ($z=$ 0.033), A2670 ($z=$ 0.076), which harbor different physical properties (see Table 1) and dispose of ancillary multifrequency data.

A85 is more massive and more luminous in X-rays than the other two systems. Despite its relaxed morphology it is known to be undergoing minor merging processes with a subcluster and groups seen in X-rays and in the optical (Durret et al., 1998; Durret, Lima Neto, & Forman, 2005; Bravo-Alfaro et al., 2009). More recently, this system has been confirmed to possess a centrally peaked profile in X-rays with a sloshing pattern (Ichinohe et al., 2015) and with some additional disturbed features (Laganá, Durret, & Lopes, 2019). The presence of jellyfish galaxies like JO201 has been associated with strong RPS in this cluster (Ramatsoku et al., 2020; Luber et al., 2022, and references therein).

A496 is known to have a very relaxed morphology (Durret et al., 2000; Boué et al., 2008a; Ulmer et al., 2011) consistent with a significant coincidence between the BCG and the cluster centroid (position and velocity) (Lopes et al., 2018). Nevertheless, Laganá, Durret, & Lopes (2019) reported disturbed features superposed onto a centrally peaked X-ray distribution. Concerning the LSS, A496 was reported as an isolated system by Chow-Martínez et al. (2014). We revisit this result in Sect. 6.

A2670 is the faintest of the three systems in X-rays, however its ICM is not negligible: RPS is blamed for strong sweep of material observed in a post-merger galaxy (Sheen et al., 2017). A2670 is the most irregular of the three systems and displays a very elongated morphology distributed along a NE-SW axis. This cluster is dynamically young, as suggested by the diffuse X-ray emission and the large offset in velocity of 333 km s^-1 (C. Caretta, priv. comm.) between the cluster and the BCG; this clearly indicates that the massive galaxy has not yet settled at the bottom of the cluster potential well.

We take advantage of Hi-data obtained with the VLA¹¹1The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. between 1994 and 1996 for A2670, and from 2001 to 2003 for A85 and A496. All observations were taken in C configuration, producing beam sizes of ${\sim}$ 24^′′$\times$ 17^′′. Table 2 gives the main observing parameters: the total integration time, the VLA configuration, the beam size, the bandwidth and the correlator mode applied. The correlator mode defines the number of polarizations applied. The correlator mode 1 uses one polarization and allows a larger number of velocity channels; mode 2 applies two polarizations and is able to produce an improvement in sensitivity by a factor of $\sqrt{2}$. Hanning smoothing was used on the data of A85 and A496, obtaining a set of 31 independent channels. For A2670 no Hanning smoothing was applied, obtaining a set of 63 channels (cubes NE and SW). The general observing strategy is based on a combination of several observing pointings at different positions and central velocities. This allowed to cover the full cluster volume beyond one virial radius and more than three times the velocity dispersion. Fig. 1 shows the total observed fields for each cluster and Table 3 gives the parameters of the final Hi data cubes. Throughout this paper we list the Hi heliocentric velocities using the optical definition.

We applied different observing strategies to each cluster depending on their redshifts and velocity dispersions. Six VLA fields were pointed side by side across A85. The individual fields have slight overlap in RA, Dec in order to avoid any gap in between. These fields have the same central velocity giving a homogeneous coverage across an area of ${\sim}$100^′ side after a mosaicing procedure (Taylor, Carilli, & Perley, 1999). The same six fields were observed three times, each one centered at a different velocity: 15,266 km s^-1, 16,500 km s^-1, 17,744 km s^-1 (see Table 3). The central velocities were chosen to slightly overlap producing a homogeneous coverage of some 3800 km s^-1, i.e. more than 3.5 times the velocity dispersion of A85.

A similar procedure was applied to A496. Since it is located at two thirds the distance of A85 the total integration time for A496 was significantly reduced. On the other hand this cluster is spread across a larger region of the sky and we had to apply ten VLA pointings with different RA, Dec positions in order to reach the one virial radius coverage. These fields went through a mosaicing procedure delivering a field of view of some 100^′. Having a low velocity dispersion (737 km s^-1), the ten fields were observed two times, each one centered at a different velocity: 9,285 km s^-1 and 10,429 km s^-1. This delivered a coverage of ${\sim}$ 2400  km s^-1, equivalent to more than 3.5 times the cluster velocity dispersion (Table 3).

A different observing strategy was applied to A2670. Located at a larger distance than the previous clusters, A2670 needed fewer VLA pointing fields to cover the virial radius. However, due to its distance, longer observing time was needed in order to reach a similar detection limit than the other clusters. The one virial radius coverage was ensured by pointing three VLA fields spread along a NE-SW axis and having some overlap in order to avoid any gaps. Due to observing time limitations, the covered area is elongated, giving priority to the NE-SW axis. Mosaicing was not applied in this cluster and the three data cubes were analyzed separately. The velocity coverage was ensured by applying a bandwidth of 12.5 MHz and correlator modes 1A for the central cube, and 1D for the NE and SW cubes. This strategy delivered a central data cube with a coverage of 4,000 km s^-1 and nearly 3,000 km s^-1 for the NE and SW cubes (Table 3).

Standard VLA calibration and imaging were carried out by using the NRAO Astronomical Image Processing System (AIPS). All final data cubes for A85 and A496 have 31 channels with velocity widths between 43.8 km s^-1 and 46.9 km s^-1. The central field A2670-C has 88 channels while 2670-NE and 2670-SW have 63 channels. All data cubes for this cluster have a channel width close to 48 km s^-1. The imaging was made by applying ROBUST$=$ 1.0 (Briggs, 1995), which is the best compromise between uniform and natural weighting. This maximizes the sensitivity while keeping high spatial resolution, delivering a final beam size of about 24^′′$\times$ 17^′′ for the three clusters. The typical rms is ${\sim}$ 0.24 mJy beam^-1 for A85 and A496, while ${\sim}$0.12 mJy beam^-1 was reached for A2670 (see Table 3). This allowed to obtain Hi-mass detection limits of the same order of magnitude for the three systems.

The search for detections and the Hi analysis were carried out using AIPS. The moment maps were produced with the AIPS task MOMNT, applied with Gaussian (space coordinates) and Hanning smoothing (velocity) to maximize the signal-to-noise ratio. We produced Hi-maps and velocity fields using a cutoff level between 1.5 and 2.0 times the rms measured in the corresponding data cube. As expected, variations of noise were found at the edge of the data cubes, in RA, Dec and velocity. We normally discarded detections found within those zones.

We obtained optical images with the 3.6m Canada France Hawaii Telescope and the MegaCam instrument, each field having a size of around $1\times 1$ deg². The south tip and the infalling filament of A85 were observed in October 2004 (program 04BF02); we exploit here the $u$ $g$ $r$ $i$ band images, with surface magnitude limits of 26.7, 27.3, 26.4, 25.7 mag arcsec^-2, respectively. Details can be found in Boué et al. (2008b). We later retrieved archive images, this time centered on A85 in the $g$ and $r$ bands. The corresponding depths are 26.3 and 26.1 mag arcsec^-2. We assembled the individual images (five in each band) to obtain frames with total exposure times of 3960 s and 8380 s.

A496 was observed with the same telescope and instrument (November 2003, program 03BF12) obtaining $u$ $g$ $r$ $i$ band images. The total exposure time is 7853 s for each image and the depths are 27.4, 27.2, 26.2, 25.6 mag arcsec^-2, respectively. Details on these data can be found in Boué et al. (2008a).

The data for A2670 were taken from the CFHT archive, with a total exposure time of 6160 s in the $u$ $g$ $r$ $z$ bands. The depths are 25.7, 26.8, 26.3, 24.2 mag arcsec^-2, respectively.

We use the spectroscopic redshift compilation by H. Andernach (priv. comm.) to define the member sample of the three clusters. This compilation considers as potential members all galaxies with published radial velocities inside a projected Abell radius ($R_{A}=2.1~{}h_{70}^{-1}$ Mpc), and within a range of ${\sim}$2500  km s^-1 from a preliminary estimation of the cluster velocity. Several observational limitations could prevent the redshift catalogues to be fully complete within the studied regions; nevertheless it remains the most reliable option with the data currently available in the literature. Taking this redshift sample, the membership catalogues are constructed by tracing caustic curves representing the escape velocity of the system in a projected phase-space diagram (see, Serra et al., 2011, and references therein for more details). The curves are constructed following Chow-Martínez (2019). All galaxies within the defined limits are taken as cluster members. The caustics applied here are intentionally relaxed (an enclosing fit) as the catalogues are intended to retain as many potential member galaxies as possible. The membership catalogues contain 616, 368, 308 objects in A85, A496, A2670, respectively. Due to its larger redshift A2670 is sampled through a larger linear radius. This bias will be taken into account when we compare the observed clusters with each other in Sect. 6.

Hi blind surveys covering large cluster volumes allow to study the gas component of all the member galaxies contained within the observed data cubes. Even those spirals not detected in Hi are very important as they trace the distribution of Hi swept spirals throughout the studied cluster. With this in mind we built catalogues (RA, Dec, $v_{rad}$) of the brightest spirals contained within the VLA data cubes. As a first step, magnitudes $g,r$ were obtained for all the member galaxies by applying SExtractor²²2https://www.astromatic.net/software/sextractor/ to the CFHT images referred above. A few zones observed by the VLA but not covered by the CFHT (Fig. 1) were similarly analyzed by using Pan-STARRS³³3https://panstarrs.stsci.edu/ frames (Flewelling et al., 2020).

The strategy to define the bright spiral catalog is as follows. We applied a limit in apparent magnitude of $g=$  19.0, 18.0, 19.5, for A85, A496, A2670, respectively. Considering the distance to each cluster, these values roughly correspond to the same absolute magnitude, M_g${\sim}$ -18.0. In r-band this limit is close to M${}_{r}=-$18.2. We convert these magnitudes to stellar mass following Mahajan et al. (2018), showing that our catalogues of bright spirals are complete in mass above a value of log$(M_{*}/M_{\odot})=9.0$. We did not apply a color-index threshold in order to include any possible red spiral. This method delivered bright objects whose morphology can be defined through visual inspection; this was done by using combined (RGB) optical images from CFHT and Pan-STARRS. All objects in the blue cloud (see next subsection) and having stellar mass above log$(M_{*}/M_{\odot})=9.0$ were considered as candidates to be spiral. Our visual inspection confirmed spiral features for most of those objects. The remaining ones, before being included in the sample of bright spirals, should present a combination of features based on a disky shape, blue ($g-r$) color, stellar mass, Hi content, and the presence of dust lanes (for the near edge-on ones). As shown below, a few red spirals are reported in the three clusters. Their belonging to the bright spiral sample was conditioned by the following criteria: having a mass above log$(M_{*}/M_{\odot})=9.0$, being detected in Hi or, alternatively, showing spiral features through our visual inspection. This output was cross-matched with the list of member objects (Sect. 2.4) in order to ensure that these spirals are cluster members. The final catalogues of bright spirals consist of 65, 46, 61 member objects in A85, A496, A2670, respectively. Hereafter we denote these objects as B-Sp.

Tracing color-magnitude diagrams (CMD) in a cluster helps to visualize the distribution of different types of galaxies populating the system. Fig. 2 shows the CMD of the three clusters, giving the color index as a function of the absolute magnitude M_r (bottom axis) and stellar mass (upper axis). The red dotted line draws the border between the red sequence and the green valley; it is defined as: ($g-r$) = 0.63 $-$ 0.02 (M_r + 20) (Masters et al., 2010).

Member galaxies in Fig. 2 are shown with grey crosses. The bright spirals defined in the previous section

are cluster members too; they are indicated with solid circles, either red or blue, indicating their position in the red sequence or in the blue cloud. Cluster members in the red sequence being brighter than M${}_{r}=-$18.2 ($i.e.$ mass above log$(M_{*}/M_{\odot})=9.0$), and having no solid circle, correspond to early types.

Galaxies above the same mass limit and lying in the blue cloud are confirmed to be part of the B-Sp sample (see previous subsection). The Hi detections are shown with blue open squares (normal objects) or with black diamonds (Hi-abnormal objects, defined in the next section). A few Hi detections are not part of the compilation of member galaxies; either no redshift was available or they are outside the velocity window defined by the membership catalog of Sect. 2.4. We added these objects as new cluster members.

Galaxies located in the blue cloud of these clusters are good candidates to be in a pre-quenching stage. Optical colors are not the best indicators for current star-forming activity due to possible differences between the time scales for quenching and for the change in color due to the stop of star formation. Nevertheless, not having optical spectra for all the objects under study, we consider that a blue color combined with a normal gas content, constitute a reasonable indicator for galaxies in a pre-quenching phase. On the other side, objects above the red line of Fig. 2 are not expected to have any intense star formation activity, therefore red circles in this figure constitute very likely passive spirals. A K-correction was applied to the three clusters. The largest K-correction values in color-index are slightly above 0.1 mag. As expected, these values correspond to red objects in A2670 (we used the K-correction calculator from GAISh⁴⁴4http://kcor.sai.msu.ru). As this correction is associated with the galaxy distance, the y-axis of Fig. 2 shows the difference in absolute magnitudes after the K-correction was applied. Galactic extinction is negligible for these clusters and no correction was applied.

The catalogues of bright spirals (B-Sp) described above provide positions and velocities that were used to seek for their Hi-emission in the corresponding data cubes. Most of these B-Sp were effectively detected in Hi but, in all three clusters, significant fractions of these spirals were not detected. These objects constitute our sample of Hi non detected spirals. This sample is complete in mass above the limit log$(M_{*}/M_{\odot})=9.0$. The search for 21 cm emission was completed by a thorough visual inspection of the Hi  data cubes. This was done blindly, disregarding all information on optical positions and redshifts of the spiral galaxies. Our search produced a list of emissions, most of them matching in position and velocity with a B-Sp. However, several candidate detections had no counterpart in the B-Sp catalog. A subsequent search upon the optical-RGB frames and redshift catalogues allowed the finding of the optical counterparts. As expected, these are blue, irregular galaxies, fainter than the limit defining the B-Sp sample.

Our strategy to search for Hi emitting galaxies includes a conservative detection threshold above 6 times the rms, delivering 10, 58, 38 Hi-detections in A85, A496, A2670, respectively. A few Hi objects are just marginally resolved by the VLA primary beam: two in A85, six in A496, eleven in A2670. All maps and velocity fields are available as online material. We used the Hi mass detection limits given in Table 3 to ensure that the sample of bright spirals (Sect. 2.5) is potentially detected with this survey. Obviously, more massive spirals have higher chances to be detected than less massive ones. Therefore, we inspect the performance of our survey detecting galaxies with different stellar masses. We displayed our Hi mass detection threshold onto a M_HI vs M_∗ plot (not shown in this work). We added the scaling relation obtained by Dénes et al. (2014) and we found that normal galaxies can be effectively detected down to log$(M_{*}/M_{\odot})=9.0$. The deeper data cubes in A496 pushed this limit down to log$(M_{*}/M_{\odot})=8.7$. This confirms that our full sample of B-Sp can be detected with the present survey under unperturbed conditions.

Tables 7, 8 and 9 give the main optical and Hi-parameters of the detected galaxies in A85, A496, A2670 respectively. The lists of non detected spirals are given in Tables 10, 11, 12 for the same clusters. Five Hi detections in A496 could not be properly measured as we did not recover their full emission; this is due to limited velocity coverage of the corresponding data cube. We consider these galaxies as Hi detections because the recovered emission has large S/N; they are indicated with black squares in Figs. 2 and 3.

In Tables 7, 8, 9, the first five columns give names, coordinates and morphological type; these data are taken from the literature (NED⁵⁵5https://ned.ipac.caltech.edu/ and HyperLEDA⁶⁶6http://leda.univ-lyon1.fr/). Column 6 and 7 give optical magnitudes and ($g-r$) colors obtained from CFHT and from Pan-STARRS images. In Column (8) we provide the stellar masses obtained from our M_r magnitudes. Columns (9) and (10) give the optical and Hi velocities, the latter in heliocentric frame. We consider the value of the central channel displaying emission as the Hi velocity. We give the Hi velocity width in Column (11) defined by the full range of channels containing emission; this constitutes a proxy for $W_{20}$. The Hi mass (Column 12) in solar masses is calculated by using the equation:

Galaxies having at least one of the four criteria mentioned above are defined as Hi abnormal. We used this sub-sample of perturbed objects to carry out a search for a correlation between Hi-deficiency and Hi-asymmetry, finding no obvious trend. However this result based on gas deficient, yet Hi detected galaxies, must be taken with caution as these objects are caught at an early stripping phase. We have no information on the 2-D gas distribution of the even more stripped spirals, $i.e.$ those not detected in Hi. Therefore a deficiency–asymmetry correlation, developing during later gas stripping stages, should not be discarded a priori.

Some of our main results are displayed in Fig. 3, Fig. 4 and in Table 4. The (RA, Dec) distribution of member galaxies in the three studied clusters is shown in Fig. 3. We use the same symbols as in Fig. 2 in order to emphasize the position of the bright spirals (blue/red circles) and the Hi-detected galaxies (open blue squares/black diamonds). The total fields covered with the VLA are drawn with blue polygons and the corresponding R₂₀₀ are indicated with red circles.

Fig. 4 shows the moment-zero Hi-maps of the detected galaxies in each cluster. Galaxies are set at their proper location and are magnified by a factor of five for clarity. We use a color code in order to separate these galaxies in three radial velocity regimes: lower, similar and larger than the cluster velocity. We give statistical significance to the velocity slots by applying the same criteria as Wang et al. (2021), given as a function of the velocity dispersion of the cluster ($\sigma_{cl}$) given in Table 1: blue color for $\Delta v_{rad}<-0.5\sigma_{cl}$, green for $-0.5\sigma_{cl}<\Delta v_{rad}<0.5\sigma_{cl}$, and red for $0.5\sigma_{cl}<\Delta v_{rad}$. The X-ray emission is shown in grey scale taking advantage of the homogeneous and large field map given by ROSAT.

Table 4 reports the total number of objects having different observational properties and lying in specific cluster regions. Data are displayed as follows: Row (1) gives the full number of cluster members, including the new members derived from the Hi velocities. Row (2) provides the number of galaxies within the VLA FoV. Rows (3) and (4) give the total number of Hi detections and bright spirals. Rows (5), (6), (7), give the number of B-Sp being Hi detected, the B-Sp non-detected and the B-Sp lying in the red sequence, respectively. Rows (8), (9), (10), give the number of galaxies detected in Hi lying in the red sequence, the Hi abnormal, and those below the limit log$(M_{*}/M_{\odot})=9.0$, respectively. The last three rows, (11), (12), (13), provide the number of galaxies projected within the virial region (Sect. 5.2) being: cluster members, bright spirals and Hi detections, respectively. The data given in Table 4 are not well suited for a direct comparison of the clusters with each other because of their different redshifts and VLA coverage. We carry out a comparison under homogeneous conditions in Sect. 6.

Fig. 3 (upper panel) shows the distribution of 616 member galaxies in A85 where the solid circles indicate the 65 B-Sp. The blue open squares and one black diamond indicate the ten Hi detected objects. The cluster members (grey crosses) are rather homogeneously distributed across A85, while the B-Sp and the Hi-detected galaxies show a very asymmetric distributions across the cluster. The SE zone has more bright spirals than all the other quadrants taken together; all the Hi-detections are projected as well onto the SE-NE regions (Fig. 3 and Fig. 4). The gas rich galaxies are projected 1.5$-$3.0 Mpc away from the cluster center, while a remarkable absence of Hi detections at lower radius suggests the presence of a harsh environment. This fact is independently supported by the presence of strongly perturbed galaxies (Venkatapathy et al., 2017; Ramatsoku et al., 2020; Luber et al., 2022).

The sample of 65 B-Sp in A85 is unexpectedly high. A significant fraction of these objects are projected around the cluster core and in the transition zone between the center and the SE (Fig. 3). Six B-Sp are reaching the red sequence (red solid circles in Fig. 2) and none of them is detected in Hi, indicating that they have gone through an advanced stripping process. On the contrary, all the galaxies detected in Hi show very blue colors indicating that they have not gone through any stripping phase.

The asymmetry in the distribution of Hi-detections is even more striking in Fig. 4, showing that all –but one– detections are in the low velocity domain, with a narrow $\Delta v$ of ${\sim}$ 15,000$-$15,200  km s^-1 (see Table 7). Their low velocity dispersion and the large difference compared with the systemic velocity of A85 (${\sim}$ 1,500  km s^-1) suggest that they could share a common origin.

The middle panel of Fig. 3 shows the distribution of 383 member galaxies in A496, with 46 B-Sp and 58 galaxies detected in Hi (symbols are the same as in Fig. 2). A first remarkable difference of A496 when compared with A85 is given by the lower number of bright spirals (46 in A496, and 65 in A85) but a much larger number of Hi-detections (58 in A496 for only 10 in A85; see Table 4). Interestingly, ${\sim}$50% of the Hi detected galaxies in A496 correspond to gas-rich, low mass objects which are not part of the B-Sp sample, $i.e.$ their stellar mass is below log$(M_{*}/M_{\odot})=9.0$. This fraction is partly explained by A496 being closer than the other studied clusters. However, we will show that some peculiarities remain after considering the observational bias due to the difference in redshift.

The distribution of Hi-detected objects in A496 (Fig. 3) is striking, with many Hi normal galaxies projected onto the central regions. The middle panel of Fig. 4 helps to solve the paradox: the objects with low velocity relative to the cluster (maps in green, in the velocity range $-0.5\sigma_{cl}<\Delta v_{rad}<0.5\sigma_{cl}$) could be dominated by a movement on the plane of the sky, with circular orbits still far from the densest ICM regions (Dressler, 1986). The sample of red maps projected close to the cluster core in Fig. 4 have large components of velocity along the line of sight. They must be at an early stage of infalling, accounting for their normal gas content.

Seventeen B-Sp in A496 are found in the red sequence; eleven of them still retain enough gas to be detected, some are classified as Hi abnormal and are indicated with black diamonds in Fig. 2 and Fig. 3. This strongly suggests that they are under an advanced process of stripping. Most of the red spirals in A496 are projected in the NW and SE quadrants; in Sect. 6 we discuss a possible relation between their gas properties and their location within the cluster.

The bottom panel of Fig. 3 shows the distribution of 315 member galaxies in A2670, emphasising the position of 61 bright spirals and 38 Hi-detections (same symbols as previous figures). From the 61 B-Sp reported in this cluster, 28 are not detected in Hi, suggesting an active sweeping process. Fig. 2 shows 8 B-Sp in the red sequence, some of them retaining enough gas to be detected in Hi.

Fig. 3 unveils a complex distribution of the Hi detections in A2670 which are mostly projected onto the cluster core and to the SW. Most of the Hi non-detected bright spirals in A2670 (solid blue/red circles without square or diamond in Fig. 3), are located near the core and to the NE. Fig. 4 shows a trend of abundant Hi rich galaxies in the SW while the NE beyond 1 R₂₀₀ is very poor in Hi. The center shows a balance between shrunk and extended Hi maps but we remind that strong projection effects might be present in this zone.

We applied an improved version of the well known $\Delta$-test (Dressler & Shectman, 1988), which is based on the parameter-$\delta$, defined by:

The $\Delta$-test detects a substructure as a concentration of galaxies showing high values of the parameter-$\delta$; this has proven to be a reliable tracer of substructures. However, this method loses some accuracy in certain cases, for instance, when two groups at different distances from the observer are projected along the same LOS. In this case the $\Delta$-test will produce one single overestimation of the parameter-$\delta$, detecting only one structure. This is critical in the absence of a dominating large structure, like in bimodal/multimodal systems. In order to improve this method we introduce dividing surfaces on a 3D distribution of RA, Dec and $v_{loc}$. This separates the sub-groups of galaxies that share the same kinematic properties within their particular caustics, but somehow detached from the kinematics of the whole cluster. That is, we do not take into account only the parameter-$\delta$, but also the $v_{loc}$ and the phase-space distribution of galaxies that are candidates to be members of the substructures (more details in Caretta et al., 2022, in prep.). Values of the parameter-$\delta$ are illustrated in Fig. 5 by black circles; these circles grow with the probability of a galaxy to belong to a group. Individual values of the parameter-$\delta$ will be available upon request.

We revisit the analysis of substructures of A85 (Bravo-Alfaro et al., 2009) by taking advantage that the number of redshifts for member galaxies has grown significantly in the last years (e.g., Boué et al., 2008b). For the full sample of cluster members we report a velocity dispersion of $\sigma_{v}=1,054$ km s^-1. We report the following substructures.

We detected a main cluster body with 330 galaxies and a $\sigma_{v}=674$ km s^-1. We clearly unveiled one single substructure located ${\sim}$1^∘ (2.5 Mpc) NW from the cluster center, with 38 members (above 10% of the total mass), $v$ = 9,338  km s^-1, and $\sigma_{v}=457$ km s^-1 (see Fig. 5). Our Hi survey covers only a small portion of that group, delivering only limited information on the cold gas associated with this structure.

We obtained the following results for A2670:
- The main cluster body emerges with 287 members, $v$ = 22,806  km s^-1, and $\sigma_{v}=805$ km s^-1.
- We unveiled two low-mass structures, the first one projected close to the center (NC, ${\sim}$10^′ to the north) with 7 members, $v$ = 22,883  km s^-1, and $\sigma_{v}=20$ km s^-1. This velocity suggests a merger with the main cluster body.
- The NE group lies ${\sim}$50^′ (4.5 Mpc), NE from the cluster center, with 14 members, $v$ = 23,115  km s^-1, and $\sigma_{v}=424$ km s^-1 (Fig. 5). This structure is just north from the VLA FoV (Fig. 3) and we do not dispose of Hi information for these galaxies.

In parallel to the $\Delta$-test described above, we applied the Serna-Gerbal (S-G) method (Serna & Gerbal, 1996) to the three clusters. This test introduces a fundamental, independent element to detect substructures: the relative potential energies of each pair of galaxies by using individual galaxy brightness as a proxy for the stellar mass. We use the $r$-band magnitude for this goal. In practice the S–G test requires an input catalogue of positions, velocities, magnitudes (converted to masses assuming a value of M/L), and a defined number $n$ of galaxies needed to build a group. We used $n=10$ as a representative value. The S–G test works hierarchically, seeking for lower level substructures within a primary (more massive) detected group. This delivers an output list of galaxies belonging to each substructure and provides the mass of the corresponding group. Absolute masses of individual structures are not highly accurate, but the relative masses of the groups, normalized to the total member sample, are very reliable.

The S–G program computed for A85 a mass of M${}_{\mathrm{tot}}=6.51\times 10^{14}$ M_⊙ for the entire cluster, and a velocity dispersion of $\sigma_{tot}=1,037$ km s^-1. The S–G test detected the main cluster body (identified as #1) with 587 galaxies, a relative mass M${}_{\mathrm{1}}=0.792\,$M_tot, and $\sigma_{1}=949$ km s^-1. A secondary group (#2) is detached from the main cluster body, lying south from the center and having 21 galaxies, M${}_{\mathrm{2}}=0.079\,$M_tot, and $\sigma_{2}=184$ km s^-1. This structure is coincident with the SB group detected by the $\Delta$-test.

In the next step we took Structure #1 as the starting point, and we unveiled two further substructures; the first one (identified as #1.1) has 144 galaxies with M${}_{\mathrm{1.1}}=0.218\,$M_tot, and $\sigma_{1.1}=527$ km s^-1. Some galaxies from this group account for the SE subcluster, detected as well with the $\Delta$-test. The second substructure (#1.2) has 14 galaxies, a relative mass M${}_{\mathrm{1.2}}=0.01\,$M_tot, and $\sigma_{1.2}=180$ km s^-1. Finally, the group #1.1 harbors two additional substructures: the first (#1.1.1) has 74 galaxies with M${}_{\mathrm{1.1.1}}=0.125\,$M_tot, and $\sigma_{1.1.1}=481$ km s^-1. This group coincides in position with the minor SW structure obtained with the $\Delta$-test. Another dozen galaxies from the structure #1.1.1 emerges as a group lying between the SB and the Filament, unveiling a more complex region, south of the core, than what is shown by the $\Delta$-test. A final subgroup (#1.1.2) appears with 10 galaxies, a relative mass of M${}_{\mathrm{1.1.2}}=0.009\,$M_tot and $\sigma_{1.1.2}=173$ km s^-1. This substructure matches with the Filament group detected with the $\Delta$-test. With the exception of the NE group, the S–G test successfully confirmed the substructures found with the $\Delta$-test in A85.

We analyzed the full sample of member galaxies but we found no major substructures in A496. Even modifying the parameter n (for instance with n=5) we did not find significantly different results. The S–G method finds that A496 is a relaxed cluster with negligible substructuring, in agreement with previous studies (e.g. Boué et al., 2008a).

In the first step, the S–G test detected two groups in A2670. The largest structure (#1) has 123 galaxies with a $\sigma_{v}$ comparable to that of the entire cluster, having M${}_{\mathrm{1}}=0.490\,$M_tot. The second group (#2) has 10 galaxies, a mass M${}_{\mathrm{2}}=0.016\,$M_tot, and $\sigma_{\#2}=215$ km s^-1. Group (#2) is placed ${\sim}$40^′ SW from the cluster center and was not detected by the $\Delta$-test (see Fig. 5). Very interestingly, this substructure embraces the zone where several Hi rich galaxies are reported (see bottom panel of Fig. 4). These objects have similar physical properties (position, velocity, gas richness), giving further support for this group to be real.

In a second run, the large structure (#1) is decomposed into two substructures (#1.1 and #1.2) having 83 and 11 galaxies respectively. The former contains M${}_{\mathrm{1.1}}=0.270\,$M_tot, and a velocity dispersion comparable to that of the cluster. The group #1.2 contains M${}_{\mathrm{1.2}}=0.180\,$M_tot, and $\sigma_{1.2}=433$ km s^-1. The fact that the substructures #1 and #1.1 present high $\sigma_{v}$, comparable to that of the whole cluster, reduces the statistical significance for these groups, as shown by Serna & Gerbal (1996). However, part of substructure #1.1 extends from the centre to the north, coinciding with the NC group found with the $\Delta$-test. The same occurs with substructure #1.2, coinciding with the NE group, reinforcing the existence of these two groups (Fig. 5).

The two methods, the $\Delta$-test and the Serna-Gerbal, confirm the very complex substructure system of A85, contrasting with the very low level shown by A496. The two tests show A2670 as an intermediate case with a few, not very massive substructures.

In order to estimate the RPS radial profiles we follow the theoretical definition proposed by Gunn & Gott (1972), $RPS=\rho_{\mathrm{ICM}}\times v_{\mathrm{rel}}^{2}$. Here $\rho_{\mathrm{ICM}}$ is the local density of the intracluster medium, and $v_{\mathrm{rel}}$ is the galaxy velocity relative to the ICM. We estimate the ICM density ($\rho_{\mathrm{ICM}}$) as a function of the cluster radius by applying the hydrostatic-isothermal $\beta-$model of Cavaliere & Fusco-Femiano (1976), $\rho(r)=\rho_{0}[1+(r/r_{\mathrm{c}})^{2}]^{-3\beta/2}.$ Table 5 gives the parameters of the $\beta$-model for the studied clusters: the core radius $r_{c}$, the central density $\rho_{0}$ and the $\beta$ value used to calculate the density as a function of the cluster radius. The corresponding RPS profiles are shown in Figure 6 where we display three representative velocities which are multiple values of the velocity dispersion of each cluster: 0.66$\sigma_{cl}$ in blue, 1.0$\sigma_{cl}$ in green, and 1.5$\sigma_{cl}$ in red. This adds physical meaning to the comparison of the RPS intensity among different clusters.

In order to quantify the effects produced by the RPS on individual objects it must be compared with the anchoring force per unit surface produced by the galaxy gravitational potential, given by $\Pi_{gal}=2\pi\,G\,\Sigma_{\mathrm{s}}\,\Sigma_{\mathrm{g}}$. Here, ${\Sigma_{s}}$ and ${\Sigma_{g}}$ are the surface densities of the stellar and gaseous discs, respectively. RPS must overcome the anchoring force in order to succeed the gas stripping. Several values of $\Pi_{gal}$ have been proposed, going from 2.0$\times$10^-13 N m^-2 for spirals in Virgo (Cayatte et al., 1994), to 2.1$\times$10^-12 N m^-2 for a Milky Way-like galaxy (Hernández-Fernández et al., 2014). One step further was made by Jaffé et al. (2018) who modelled the anchoring force as a function of the galaxy radius, the stellar mass and the gas fraction of a sample of spirals. We follow these authors in order to seek for values of $\Pi_{gal}$ being representative for our sample of B-Sp.

From Fig. 2 we see that most of the B-Sp are in the range of mass between log$(M_{*}/M_{\odot})=9.0$ and log$(M_{*}/M_{\odot})=10.5$. No significant difference in mass is expected between our sample and that studied by Jaffé et al. (2018), so we use their Figure 5 (and Fig. 4 in Jaffé et al., 2019), which provide the anchoring force at a radius of twice the disc scale length. As we are interested in the RPS effects exerted at the external edge of spirals, we take this value as a proxy for the anchoring force acting at the outskirts of the galaxy. This strategy delivers $\Pi_{gal}$ in the range from 10^-12 to 10^-13 N m^-2. We take the middle value 5.0$\times$10^-13 N m^-2 as a reference for our sample of B-Sp. As a comparison, the $\Pi_{gal}$ value reported by Cayatte et al. (1994) predicts very strong RPS effects, even in the low velocity and at large cluster-centric domains (see Fig. 6). The opposite is given by the value of $\Pi_{gal}$ proposed by Hernández-Fernández et al. (2014). We indicate the three values of $\Pi_{gal}$ in Figure 6. It is clear that the reference value $\Pi_{gal}=$ 5.0$\times$10^-13 N m^-2 is more realistic for the bright spirals under study so we use it along the forthcoming analysis.

Some uncertainties are introduced though. Particularly important is the dispersion in the relation between the anchoring force and the stellar mass, with a strong dependence on the galaxy radius. In this respect Jaffé et al. (2018) have shown that the anchoring force is strongly dependent on the degree of mass concentration; this implies that low mass objects with a high central density could have a similar $\Pi_{gal}$ at their center than a more massive/extended galaxy at its edge. Despite these uncertainties we will show (Sect. 6) that our reference value of $\Pi_{gal}$ match reasonably well with our observations on a statistical basis.

We repeat this analysis for the low-mass galaxies, $i.e.$ those in the range of stellar mass log$(M_{*}/M_{\odot})$ between 8.0–9.0, obtaining an upper limit for the anchoring force of $\Pi_{gal}=$ 10^-13 N m^-2. Galaxies in this low regime of $\Pi_{gal}$ are predicted to be easily stripped of their cold gas when they approach the virial radius. We further discuss about these objects in Sect. 6.

The projected phase-space (PPS) diagram of a cluster helps to distinguish virialized from infalling galaxies (e.g. Kent & Gunn, 1982; Geller et al., 2013; Oman, Hudson, & Behroozi, 2013; Jaffé et al., 2015; Rhee et al., 2017; Cortese, Catinella, & Smith, 2021). We present a version of PPS where the parameters $\Delta$v and the radius are normalized by the velocity dispersion $\sigma_{cl}$, and by R₂₀₀, respectively. On the y-axis we display the absolute value of the peculiar velocity ($|\Delta v_{rad}|/\sigma_{cl}$). On the x-axis the projected position relative to the cluster center is normalized by R₂₀₀.

Fig. 7 shows the PPS diagrams for A85, A496 and A2670 with the same symbols as in previous figures. The escape velocity of each cluster is shown with a solid grey line; these are more conservative than the caustics used to define the member sample and helps to obtain a better contrast between virialized and infalling galaxies (e.g. Jaffé et al., 2015; Wang et al., 2021). This escape velocity was estimated on the basis of a NFW halo with a concentration of $c=6$; other values of this parameter do not change the escape velocity significantly. The orange dotted line draws the boundary where RPS equals the anchoring force, $\Pi_{gal}=$ 5.0$\times$10^-13 N m^-2.

At the lower left corner of the PPS diagram we find the galaxies projected at short cluster-centric distances and having low (radial) velocities relative to the cluster. These objects are likely settled onto the cluster potential after completing –at least– one pericentric passage. This part of the diagram is identified as the virialized zone and we define its boundaries as follows. On the x-axis, we followed authors (e.g. Jaffé et al., 2015; Kopylova & Kopylov, 2018) who propose that the virial radius lies in the range 1.2–1.4 R₂₀₀. On the y-axis the criteria to define the virial zone is based on statistics, by considering a Gaussian distribution of velocities. We use the value of 1.5 $\sigma$ as a more restrictive limit (compared with 3 $\sigma$) in order to avoid contamination by foreground/background galaxies. Considering these boundaries we define the virialized zone as the triangle limited by the solid red line in Fig. 7. This triangle zone is more conservative than the corresponding squared region and has been applied by authors like Jaffé et al. (2015) and Wang et al. (2021). On the opposite corner of the PPS diagrams, galaxies infalling for the first time lie in the large radii/velocity domain. After a few $10^{9}$ yrs, equivalent to the crossing time for a typical cluster, infalling galaxies will lose kinetic energy due to dynamical friction and violent relaxation, and they will settle deep into the cluster gravitational potential.

The distribution of different types of galaxies in the PPS diagrams offer important hints on the dynamical stage of the three clusters. The top panel of Fig. 7 shows A85 behaving as a typical evolved system with a high concentration of member galaxies within the virialized zone. The escape velocity curve in this plot is lower than the one reported by Bellhouse et al. (2017); this is due to the higher mass used for A85 by these authors $\text{M}_{200}=1.58\times 10^{15}\text{M}_{\odot}$ We see a clear trend where, with one single exception, all B-Sp lying below the escape velocity curve are not detected in Hi, suggesting that these galaxies had gone through an advanced process of gas stripping. The gas rich objects in A85 lie in the high radius, high velocity domain, indicating that they approach the cluster for the first time. Despite some similarities between A85 and other systems like A693 ($e.g.$ mass, morphology and dynamical stage), their distribution of galaxies across their PPS differs significantly (Jaffé et al., 2015).

Similar to A85, A496 harbors a crowded virialized zone (medium panel of Fig. 7). This is in agreement with the advanced level of relaxation reported by different authors and with the low degree of substructuring in A496 reported in this work (Sect. 4). The striking number of Hi rich galaxies lying within the virialized zone can be explained by projection effects as described in Sect. 3.2. The PPS diagram of A496 shows a slight trend of Hi abnormal galaxies (black diamonds) and red spirals (red solid circles) being more abundant towards the virialized zone.

A2670 is dynamically young accounting for the complex distribution of galaxies in the PPS diagram and a virialized zone significantly less concentrated than in the other two systems. In this respect A2670 looks like other young clusters like Hydra (Wang et al., 2021). A number of Hi detections is projected onto the virialized zone, certainly explained by projection effects. Other Hi galaxies are located in the high radius ($\geq 1.5\,r$/R₂₀₀) domain, matching with the substructure reported onto the SW in Sect. 4. Several of these detections correspond to abnormal Hi galaxies (black diamonds) and to red spirals which are good candidates to be under pre-processing.

The results described so far provide undeniable evidence on the important role played by environment in the transformation of spirals in the three studied clusters. The global cluster Hi deficiency is an important parameter that correlates with the properties of galaxies in most of the nearby clusters. In this work we approach the cluster deficiency as the ratio of high Hi deficient B-Sp (i.e. those not detected in Hi) normalized by the total number of B-Sp found within each system. Cluster deficiencies calculated in this way are: 0.86, 0.37, 0.46, for A85, A496, A2670, respectively. A85 appears as an extreme case of Hi deficiency and, since it is the most massive of the three studied systems, this result matches the model prediction of more massive halos producing larger gas loss (e.g. Donnari et al., 2021).

Next we discuss our results in the frame of four milestone results and predictions on galaxy evolution in clusters which are based on observations and on theoretical models: (1) The dependency of individual galaxy properties (like gas content and SF quenching) on the cluster-centric distance. (2) The commonly accepted role played by RPS in galaxy transformation, at least in the inner cluster regions. (3) The pre-processing of galaxies at large cluster-centric radii is thought to be responsible for the transformation of large fractions of galaxies. (4) Models of galaxy evolution predict the transformation (quenching and gas sweeping) of large fractions of satellite galaxies previous to their infalling (we refer to the following works and references therein, Peng et al., 2010; Rhee et al., 2017; Jung et al., 2018; Roberts et al., 2019; Cortese, Catinella, & Smith, 2021; Donnari et al., 2021; Healy et al., 2021; Boselli et al., 2022).

We explore how the three studied clusters fit within these predictions taking into account the results described in the previous sections. With this aim we use our sample of B-Sp as test particles for environment effects. As a first approach we consider that a normal Hi content plus a blue color constitute an indicator of a non-quenched galaxy. As mentioned previously, optical colours are not the most accurate tracers for quenching but they are often used as an alternative, distinguishing active from non-active galaxies, for instance in statistical studies as the present one. In the following discussion we take our sample of B-Sp as analogues of the satellite galaxies defined in cosmological simulations. We observe that none of the B-Sp dominates any of the groups reported in Sect. 4, giving support to this approach. Current models (e.g. Xie et al., 2020) predict that evolution of low mass galaxies, those having log$(M_{*}/M_{\odot})<$ 10.3 (matching the range of our B-Sp) is dominated by environmental effects. Above this threshold AGN feedback is expected to drive the galaxy quenching.

Here we compare some physical properties of the studied clusters with each other. In order to avoid part of the observational bias we focus on the inner regions, more specifically on the largest physical radius being observed with the VLA in the three systems. Based on this criteria we take the conservative value of 1 R₂₀₀ as comparison radius. Obviously, some bias remain due to different cluster distances which has an effect on the population of Hi detected galaxies. For instance, typical Hi masses reported in A85 and A2670 approach 8 $\times$ 10⁹  M_⊙, while in A496 this value is ${\sim}$2 $\times$ 10⁹ M_⊙. As a consequence, in this last cluster we report more Hi galaxies in the low mass regime (Sect. 3). Nonetheless, some peculiarities described below remain even after considering this observational bias. Table 6 resumes the data corresponding to 1 R₂₀₀ for the three clusters.

These effects are well explained by RPS. We showed (Sect. 5.1) that this mechanism might exceed the anchoring force up to a radius close to 0.6 $r$/R₂₀₀ for typical ($1\sigma$) velocities, and up to 0.9 $r$/R₂₀₀ for faster infallers ($1.5\sigma$; see Fig. 6). As a matter of fact, the central regions of A85 could be an extreme case even compared with Coma, where a number of gas deficient spirals are still detected within 1 Mpc, with similar Hi mass limit detection than the present work (Bravo-Alfaro et al., 2000; Healy et al., 2021; Molnár et al., 2022). In A85 the only Hi  projected close to the center is the high-speed infalling jellyfish JO201, which shows signs of strong ram pressure stripping (Ramatsoku et al., 2020; Luber et al., 2022).

In parallel to RPS, we inspect the role played by pre-processing and by the shocked interface between A85 and the SE-group, in the transformation of infalling spirals. A large majority of the B-Sp which are not detected in Hi  are located in the SE quadrant, where the shocked interface is found. The complexity of this region is consistent with the presence of two groups (the South Blob and the Filament groups) in addition to the SE subcluster itself (Fig. 5). Another jellyfish galaxy (JO200) is located in this zone being so deeply stripped that it was not detected in Hi down to a mass limit of 9 $\times$ $10^{8}$ M_⊙ (A85[DFL98]286 in Table 10). Venkatapathy et al. (2017) already proposed this galaxy to be in a more advanced stage of transformation than JO201.

A remaining number of Hi non-detected B-Sp galaxies lies in the NW quadrant of A85 (Fig. 3). These objects are seen slightly beyond the escape velocity curve of the PPS of Fig. 7. No substructures are reported in that zone; therefore we favor the backsplash hypothesis (Gill, Knebe, & Gibson, 2005; Hirschmann et al., 2014). This suggests that these galaxies had entered A85 from the SE and now we see them projected to the NW. Timescales for RPS quenching (${\sim}$1 Gyr) and for cluster crossing (${\sim}$3 Gyr) are in good agreement. Considering that morphology transformation, disregarding the responsible mechanism, takes longer timescales (${\sim}$10 Gyr), backsplash is a plausible explanation for the spiral morphology and their location within the cluster (Boselli & Gavazzi, 2006, 2014; Cortese, Catinella, & Smith, 2021).

As a first approach to quantify pre-processing in A85 we obtained the number of Hi perturbed galaxies, and those non-detected, within each of the groups reported in Sect. 4. We obtained the following numbers: SB:3, Filament: 4, SE: 4, SW: 0, NE: 1. This suggests that pre-processing could be at work within the groups located SE of A85. However it is difficult to separate pre-processing from the effects of the interface shocking region. The later covers very well the zone where the B-Sp not detected in Hi are found.

The sample of blue, Hi normal galaxies in A85 projected between 1.5 Mpc and 3.0 Mpc east of the cluster center raises some unanswered questions. They are confirmed as cluster members and, being in the stellar mass range log$(M_{*}/M_{\odot})$ between 9.0 and 10.0, these objects are predicted to have high probability to be quenched before their complete infall onto a massive cluster like A85 (e.g. Bahé et al., 2019). For instance, Xie et al. (2020) estimates that ${\sim}$60% of galaxies would be quenched in the mass range given above and halos (clusters) with masses above log$(M_{halo}/M_{\odot})=$ 14.0, which is the case for this cluster.

In any case the gas rich galaxies in A85 seem to be part of a continuous flow of newcomers occurring during several Gyr. This scenario is supported by the fact that A85 is part of the supercluster MSCC-039 which includes eleven systems (Chow-Martínez et al., 2014). The large scale structure joining A85 with its neighbors, in particular A117A, projected at ${\sim}$12 Mpc SE of A85 (C. Caretta 2021, priv. comm.), could be responsible for the feeding of fresh Hi galaxies and with the presence of the SE subcluster.

Taking the full sample of B-Sp as test particles for environment effects we find that 37% of them have lost a large fraction of gas and were not detected in Hi (Table 4). Furthermore, among the Hi detected B-Sp a number of them are reported as abnormal (Table 8). If we consider the color index, one third of the B-Sp in A496 are located in the red sequence. All these results clearly indicate an important degree of spiral transformation in this cluster.

Another intriguing result in A496 is the large number of Hi detected galaxies with masses below log$(M_{*}/M_{\odot})$=9.0, some of them having Hi mass $\geq$ 3$\times$10⁹ M_⊙ (see Table 8). In similar clusters, like Hydra, the average Hi mass for the same range of stellar mass is significantly lower, below 10⁹ M_⊙ (Wang et al., 2021). Even considering the closer distance of A496 (the nearer the cluster, larger the fraction of low-mass Hi detected objects) we expected similar galaxies to be found in A85 and A2670. Taking into account our Hi detection limits and by scaling the corresponding Hi flux to the distance of A85 and A2670, our survey would be capable to detect such gas-rich galaxies with stellar masses around log$(M_{*}/M_{\odot})$=8.5. However these objects are absent, more conspicuously in A85.

The low-mass, gas-rich and very blue galaxies in A496, are very likely in a pre-quenching phase. These objects deserve a more detailed study; current models predict strong pre-processing and a significant quenching fraction, going up to 80% for log$(M_{*}/M_{\odot})=$ 9.0, at $z=0$ (Donnari et al., 2021). In a forthcoming paper we will try to determine, statistically, how close to a massive halo these low-mass galaxies last before extinguishing. This could help to restrict theoretical models.

Beside the projection effects, the full distribution of Hi galaxies in A496 (Fig. 3) deserves a closer look. The NE quadrant is almost empty of Hi objects. In the south and south-east regions, all the Hi detections belong to the high velocity domain (red maps in Fig. 4), suggesting a common origin. Finally the north-west is the richest zone of both, Hi detections and member galaxies. This could be associated with the only substructure found in this cluster (Sect. 4). Concerning pre-processing, the very low degree of substructures suggest that this mechanism is not playing a major role transforming infalling galaxies in A496. Instead, the cluster environment seems to dominate the processing of spirals.

We explored the surrounding large scale structure of A496 in order to better understand the replenishment of gas-rich galaxies. We analyzed 6dF Survey data within an area of 12 degrees around A496; the center of this cluster is indicated with a black solid circle in Fig. 8. We applied a galaxy request restricted to the redshift range of 0.0255 – 0.0407 (7552 – 11962  km s^-1), corresponding to the 3$\sigma$ interval of A496. The output of 498 galaxies is displayed in the same figure. This plot shows in green the positions of all the galaxies found in 6dF while the objects in red correspond to the selected redshift interval. Important galaxy concentrations are clearly seen, drawing a filament in the NW-SE direction, having redshifts of 0.037 and 0.038 (the redshift of A496 is 0.033). The nearest structure to A496 lies to the NW and could be at the origin of the large number of gas-rich objects in the NW zone of A496 and the substructure reported in the same region (Fig. 5). Other observational evidences support this filament of large scale structure, like the X-ray cluster MCXC J0445.1-1551, SE from A496. This system was not identified originally by Abell and is found at RA=71.5, Dec=-16, centered on the cD galaxy NGC 1650. These filaments account for the observed replenishment of gas rich galaxies, in particular from the NW.

We use the B-Sp properties to investigate the strength of the environment effects in A2670. Fig. 2 shows that a number of B-Sp have reached the red sequence; five of them retain enough gas to be detected in Hi. Interestingly, almost half of the Hi detected galaxies are classified as abnormal (Table 4) confirming a rather harsh cluster environment. The fraction of Hi abnormal objects within 1 R₂₀₀ is much larger in A2670 (0.48) than the other two clusters. We quantify the importance of pre-processing in A2670 by contrasting the positions of the Hi perturbed B-Sp with the substructures reported in Sect. 4. We found only two of such galaxies in the north-center (NC) group and four in the SW, including 3 Hi abnormal and one non-detected object. The SW region seems to be the only zone of this cluster where pre-processing is at work.

The elongated morphology of A2670, following a NE-SW axis, is confirmed by the distribution of member galaxies in Fig. 3. The global positions of Hi detections help to trace the accretion history of A2670. When we focus beyond 1.0 R₂₀₀, a strong asymmetry is seen between the very Hi rich south-west and the Hi poor north-east region (Fig. 4). This strongly suggest the former as the preferred direction of infalling spiral galaxies.

In order to explore the global behavior of the large sample of bright spirals (B-Sp) collected in this work we create a composite cluster constituted by the three studied systems. The left panel of Fig. 9 shows the distribution, on the plane of the sky, of all member objects and all the bright spirals (we use the same symbols than in previous plots). The right panel of Fig. 9 shows the PPS diagram of the composite cluster. In order to properly display the three systems in a single plot we normalize the galaxy cluster-centric distance by the corresponding R₂₀₀. The relative velocities are normalized by the corresponding cluster velocity dispersion $\sigma_{cl}$ (see Luber et al., 2022).

We study the behaviour of the B-Sp across the composite cluster by looking for trends between their distribution and their Hi properties. We split the full catalog of B-Sp in two sub-samples named: normal B-Sp and disturbed B-Sp. The former is constituted by those Hi detected B-Sp, displaying normal distribution and gas content. The sub-sample of disturbed B-Sp includes those being Hi abnormal plus those Hi non-detected. Defined in this way, the disturbed sub-sample includes galaxies undergoing strong degrees of transformation. Next we estimate the mean projected distances for each sub-sample and we compare them. We use units of $r$/R₂₀₀, obtaining the following results:

This confirms the expected scenario where Hi disturbed spirals are located, on average, closer to the cluster center than Hi normal ones. The difference between means is, however, not found to be significant enough ($1-p\lesssim 90\%$) by the Welch’s t-test. As we next show, this is most likely due to the heterogeneous trends seen in the individual clusters. In order to quantify such trends we contrasted the average distances of the same sub-samples within each system:

This confirms, in A85, the very strong trend where disturbed B-Sp lay at lower cluster-centric radius (confidence value $>$ 99%). Oppositely, A496 has a very atypical behaviour with an inverse trend of normal B-Sp located, in average, at lower radius than their disturbed counterparts. Despite the projection effects cited above, such a trend has never been observed in other clusters. A2670 is an intermediate case where the difference between $\mu_{1}$ and $\mu_{2}$ has low statistical significance. Other statistical methods produce similar results, for instance the Two-sample test for the mean, considering the one-tail case of the z-Test.