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The steps performed are bias subtraction, bad pixel interpolation, cosmic ray cleaning, background fitting and subtraction, tracing and extraction of the spectra, and finally wavelength calibration.
The steps performed are bias subtraction, bad pixel interpolation, cosmic ray cleaning, background fitting and subtraction, tracing and extraction of the spectra, and finally wavelength calibration.
This last step was somewhat challenging, since we did not manage to take useful arc-lamp calibration data with the available wide-slit masks.
This last step was somewhat challenging, since we did not manage to take useful arc-lamp calibration data with the available wide-slit masks.
In the end, the only way to obtain a meaningful wavelength scale was to use the spectral lines of the target and comparison star themselves, and compute the wavelength solutions for each night on the basis of the respective time-averaged combined spectra.
In the end, the only way to obtain a meaningful wavelength scale was to use the spectral lines of the target and comparison star themselves, and compute the wavelength solutions for each night on the basis of the respective time-averaged combined spectra.
Luckily, subdwarf B stars show a number of easily identifiable Balmer and Helium lines at short wavelengths, and we were thus able to obtain a reasonably accurate (within a few A)) wavelength calibration for at least the blue part of the target spectrum.
Luckily, subdwarf B stars show a number of easily identifiable Balmer and Helium lines at short wavelengths, and we were thus able to obtain a reasonably accurate (within a few ) wavelength calibration for at least the blue part of the target spectrum.
This is sufficient for our purposes, since we do not attempt to detect radial velocity variations arising from the pulsations due to the low resolution of the spectra and the arbitrary wavelength variations induced by spatial shifts of the target across the wide slit.
This is sufficient for our purposes, since we do not attempt to detect radial velocity variations arising from the pulsations due to the low resolution of the spectra and the arbitrary wavelength variations induced by spatial shifts of the target across the wide slit.
The red part of the target spectrum is largely devoid of features, and the wavelength solution obtained cannot be trusted for quantitative analysis.
The red part of the target spectrum is largely devoid of features, and the wavelength solution obtained cannot be trusted for quantitative analysis.
For the comparison star, we were able to identify several lines on both chips, and achieve a rough wavelength calibration.
For the comparison star, we were able to identify several lines on both chips, and achieve a rough wavelength calibration.
However, since these measurements are only used in wavelength-integrated form in what follows, this calibration is needed as nothing more than a guideline.
However, since these measurements are only used in wavelength-integrated form in what follows, this calibration is needed as nothing more than a guideline.
Figure 1 shows a typical individual spectrum for both the target and comparison star, taken from the first night where both chips were available.
Figure \ref{spectra} shows a typical individual spectrum for both the target and comparison star, taken from the first night where both chips were available.
Note that the comparison star data were chopped at the short-wavelength end of the red chip as this part of the spectrum was saturated.
Note that the comparison star data were chopped at the short-wavelength end of the red chip as this part of the spectrum was saturated.
The high quality is striking considering the relative faintness of the target (V~15.67from?) and the short exposure time used (25 s).
The high quality is striking considering the relative faintness of the target \citep[V$\sim$15.67 from][]{kilkenny2006} and the short exposure time used (25 s).
In the central part of the blue target spectrum, we measure a signal-to-noise S/N ~ 55, while it is lower in the red part at S/N ~ 25.
In the central part of the blue target spectrum, we measure a signal-to-noise S/N $\sim$ 55, while it is lower in the red part at S/N $\sim$ 25.
For the comparison star an accurate measurement of the S/N is complicated by the numerous metal lines, but we can roughly estimate S/N ~ 90 for the central part of the blue spectrum and S/N > 100 for the red part.
For the comparison star an accurate measurement of the S/N is complicated by the numerous metal lines, but we can roughly estimate S/N $\sim$ 90 for the central part of the blue spectrum and S/N $\geq$ 100 for the red part.
It is apparent from looking at the relative depth of Ha compared to the other Balmer lines in the target spectrum that the red part is a lot less useful than the blue part in terms of S/N of the lines, mostly due to the intrinsically low flux of sdB stars at longer wavelengths.
It is apparent from looking at the relative depth of $H\alpha$ compared to the other Balmer lines in the target spectrum that the red part is a lot less useful than the blue part in terms of S/N of the lines, mostly due to the intrinsically low flux of sdB stars at longer wavelengths.
Adding to this the fact that we were able to obtain only half a night's worth of red data, and that an accurate wavelength calibration was not possible, they become unsuitable for a detailed interpretation.
Adding to this the fact that we were able to obtain only half a night's worth of red data, and that an accurate wavelength calibration was not possible, they become unsuitable for a detailed interpretation.
Therefore, the analyses presented in the following sections are based on the blue data only.
Therefore, the analyses presented in the following sections are based on the blue data only.
The first step in the data analysis was to obtain pulsation frequencies for our data set, and at the same time assure the quality of the measurements retained for further analysis.
The first step in the data analysis was to obtain pulsation frequencies for our data set, and at the same time assure the quality of the measurements retained for further analysis.
In order to maximise the S/N we integrated the flux of each spectrum over the wavelength range of interest (3650—4950 À) and produced broadband fluxes for the
In order to maximise the S/N we integrated the flux of each spectrum over the wavelength range of interest $-$ 4950 ) and produced broadband fluxes for the
thhe Local Group. this supports the view that the local deusity is near the global onc.
he Local Group, this supports the view that the local density is near the global one.
and cover enough area to sample the faint end of the LF even at large distances from the cluster centres.
and cover enough area to sample the faint end of the LF even at large distances from the cluster centres.
With these data we can exploit HST's superior photometric performance and stability, small point spread function (especially important for dwarf galaxies), high spatial resolution, and the availability of numerous deep ‘blank’ fields (e.g., COSMOS, EGS, GOODS) to provide a homogeneous set of data (taken under the same conditions) for statistical subtraction of foreground and background galaxies lying along the cluster line of sight.
With these data we can exploit HST's superior photometric performance and stability, small point spread function (especially important for dwarf galaxies), high spatial resolution, and the availability of numerous deep `blank' fields (e.g., COSMOS, EGS, GOODS) to provide a homogeneous set of data (taken under the same conditions) for statistical subtraction of foreground and background galaxies lying along the cluster line of sight.
In Pracyetal.(2004) we applied this method to a wide R band mosaic of WFPC2 fields in Abell 2218 and derived a LF down to Mg—12+5logh, showing that the LF slope appears to steepen outwards and that the faintest dwarfs avoid the cluster centre.
In \cite{pracy04} we applied this method to a wide $R$ band mosaic of WFPC2 fields in Abell 2218 and derived a LF down to $M_R \sim -12 + 5 \log h$, showing that the LF slope appears to steepen outwards and that the faintest dwarfs avoid the cluster centre.
Harsono&DePropris(2007,2009) have derived a deep composite LF in six bands for five clusters at «z>=0.25 and find that the population of dwarf galaxies down to M;=—1445logh was already present, fully assembled and lying on the red sequence at z~0.3, but find no faint end upturn.
\cite{harsono07,harsono09} have derived a deep composite LF in six bands for five clusters at $<z>=0.25$ and find that the population of dwarf galaxies down to $M_z= -14 + 5\log h$ was already present, fully assembled and lying on the red sequence at $z \approx 0.3$, but find no faint end upturn.
Here we present a study of the faint end of the J band (F814W) LF in Abell 1689 (2=0.183) over a 10' field imaged with HST WFPC2.
Here we present a study of the faint end of the $I$ band (F814W) LF in Abell 1689 $z=0.183$ ) over a $10'$ field imaged with HST WFPC2.
The next section describes the data and photometry, while we present our results and discussion in the following sections.
The next section describes the data and photometry, while we present our results and discussion in the following sections.
We adopt the WMAP7 cosmological parameters: Όλι=0.27, O4=0.73 and Ho=71 km s! Mpc-!.
We adopt the WMAP7 cosmological parameters: $\Omega_M =0.27$, $\Omega_{\Lambda}=0.73$ and $_0=71$ km $^{-1}$ $^{-1}$.
The data used in this paper consist of a 4x4WFPC2 mosaic of Abell 1689 covering ~10’x on the sky with exposure times of 1800s in the V (F606W) band and 2300s in the 7 (F814W) band.
The data used in this paper consist of a $4 \times 4$WFPC2 mosaic of Abell 1689 covering $\sim 10' \times 10'$ on the sky with exposure times of 1800s in the $V$ (F606W) band and 2300s in the $I$ (F814W) band.
The were retrieved as processed and drizzled files from the imagesHST Legacy Archive fully(PID: 5993; PI: Kaiser).
The images were retrieved as fully processed and drizzled files from the HST Legacy Archive (PID: 5993; PI: Kaiser).
In order to determine the LF of cluster galaxies we need to estimate the contribution to the total counts in the cluster line of sight from galaxies in the galaxyfield (in the foreground or background).
In order to determine the LF of cluster galaxies we need to estimate the contribution to the total galaxy counts in the cluster line of sight from galaxies in the field (in the foreground or background).
We use the J band counts in the COSMOS field (Leauthaudetal.2007).
We use the $I$ band counts in the COSMOS field \citep{leauthaud07}.
. These counts have similar photometric depth to our data (the exposuretimes are similar, but the ACS is about a factor of 2 more efficient than WFPC2), cover a large area (1.64 deg?; therefore minimizing the effects of cosmic variance) and are taken in a closely related filter.
These counts have similar photometric depth to our data (the exposuretimes are similar, but the ACS is about a factor of 2 more efficient than WFPC2), cover a large area (1.64 $^2$; therefore minimizing the effects of cosmic variance) and are taken in a closely related filter.
We therefore expect that we can use these counts to decontaminate our dataset statistically and recover the LF of cluster members.
We therefore expect that we can use these counts to decontaminate our dataset statistically and recover the LF of cluster members.
For we analyze our data in the same manner as Leauthaudconsistency,etal.(2007):: we run the Sextractor (Bertin&Arnouts1996) package twice and with the same parameters as used for the COSMOS field.
For consistency, we analyze our data in the same manner as \cite{leauthaud07}: we run the Sextractor \citep{bertin96} package twice and with the same parameters as used for the COSMOS field.
A first pass with coarse search parameters is used to detect the galaxies without deblending them, while a second pass brightwith finer search parameters is used for the faint galaxies.
A first pass with coarse search parameters is used to detect the bright galaxies without deblending them, while a second pass with finer search parameters is used for the faint galaxies.
All detections were visually inspected to remove spurious sources, artifacts and especially arclets.
All detections were visually inspected to remove spurious sources, artifacts and especially arclets.
All photometry was calibrated to the AB system using published zeropoints.
All photometry was calibrated to the AB system using published zeropoints.
By this approach we are able to use the COSMOS counts for our background removal.
By this approach we are able to use the COSMOS counts for our background removal.
In addition, we also measured two aperture magnitudes in V and J (in an aperture equivalent to 5 -! kpc) in order to determine the galaxy colors, identify the red sequence, and use this to estimate the red sequence luminosity function.
In addition, we also measured two aperture magnitudes in $V$ and $I$ (in an aperture equivalent to 5 $h^{-1}$ kpc) in order to determine the galaxy colors, identify the red sequence, and use this to estimate the red sequence luminosity function.
We also measure the and of galaxies: these are used for a ellipticitycompanion paper on positionthe anglealignment effect (Hung et al.
We also measure the ellipticity and position angle of galaxies: these are used for a companion paper on the alignment effect (Hung et al.
2010, in preparation).
2010, in preparation).
Star-galaxy separation is carried out using the Umaz vs. T diagram shown in Figure 1 (Leauthaudetal.2007),, where µπιαα 15 the central surface brightness of each object.
Star-galaxy separation is carried out using the $\mu_{max}$ vs. $I$ diagram shown in Figure 1 \citep{leauthaud07}, where $\mu_{max}$ is the central surface brightness of each object.
Stars will define a tight sequence in this plot, while galaxies will occupy a cloud of points at lower [imax for a given total luminosity.
Stars will define a tight sequence in this plot, while galaxies will occupy a cloud of points at lower $\mu_{max}$ for a given total luminosity.
The adopted discriminant between stars and galaxies is shown in Figure1.
The adopted discriminant between stars and galaxies is shown in Figure1.
As a check on our method, we determined galaxy counts in the J band for the two Hubble Deep Fields (Williamsetal.1996;Casertano2000) and verified that we obtain good agreement with the COSMOS counts reported by Leauthaudetal.(2007).
As a check on our method, we determined galaxy counts in the $I$ band for the two Hubble Deep Fields \citep{williams96,casertano00} and verified that we obtain good agreement with the COSMOS counts reported by \cite{leauthaud07}.
. In Figure 2(a) we show the LF of galaxies in Abell 1689 for the entire area covered by the WFPC2 observations.
In Figure 2(a) we show the LF of galaxies in Abell 1689 for the entire area covered by the WFPC2 observations.
We subtracted the scaled fore/back-ground counts derived from the COSMOS field, assuming Poissonian errors for the galaxy counts and including terms due to clustering errors (as per Huangetal.1997;Driver 2003)).
We subtracted the scaled fore/back-ground counts derived from the COSMOS field, assuming Poissonian errors for the galaxy counts and including terms due to clustering errors (as per \citealt{huang97,driver03}) ).
All errors are added in quadrature.
All errors are added in quadrature.
It is clear that the data are not a good fit to a standard Schechter function: the LF appears to flatten at intermediate and presents a rise at faint luminosities.
It is clear that the data are not a good fit to a standard Schechter function: the LF appears to flatten at intermediate magnitudes and presents a steep rise at faint luminosities.
This is magnitudessimilar to the local deep steepcomposite LFs for Sloan and RASS clusters observed by Popesso and to the original claims for a steep faint end upturn of the LF in clusters of galaxies (Driveretal.1994;DePropris 1995)..
This is similar to the local deep composite LFs for Sloan and RASS clusters observed by \cite{popesso06} and to the original claims for a steep faint end upturn of the LF in clusters of galaxies \citep{driver94,depropris95}. .
Following Popessoetal. (2006), we fit our data with a combination of a Schechter function and a power law: where ®*, M* and a are the usual Schechter function
Following \cite{popesso06}, , we fit our data with a combination of a Schechter function and a power law: where $\Phi^*$ , $M^*$ and $\alpha$ are the usual Schechter function
onLipparcos data. indicate that there is no apparent slope in the solar neighbourhood AMI for stars vounger than 10 Cyr (Carraro et al.
on data, indicate that there is no apparent slope in the solar neighbourhood AMR for stars younger than 10 Gyr (Carraro et al.
1998: Ne Bertelli 1998): the slope determined for the older stars agrees very well with Twarog's (1980) original determination.
1998; Ng Bertelli 1998); the slope determined for the older stars agrees very well with Twarog's (1980) original determination.
The most striking feature of these AMI determinations is the enormous intrinsic scatter around the average relation. which is generally obtained by binning in age intervals.
The most striking feature of these AMR determinations is the enormous intrinsic scatter around the average relation, which is generally obtained by binning in age intervals.
This scatter. for a given age comparable to the overall increase in metallicity over the lifetime of the Galaxy. makes the correlation for the vounger stars rather weak.
This scatter, for a given age comparable to the overall increase in metallicity over the lifetime of the Galaxy, makes the correlation for the younger stars rather weak.
Although it has been found to be relatively easy to explain the average trend by theoretical models involving infall of eas. either from the intergalactic medium or from the Galactic thick disc (Somuner-Larsen Yoshii 1990). at à rate of about half the star formation rate (SER. c.g. Pwarog 1980: Carlherg ct al.
Although it has been found to be relatively easy to explain the average trend by theoretical models involving infall of gas, either from the intergalactic medium or from the Galactic thick disc (Sommer-Larsen Yoshii 1990), at a rate of about half the star formation rate (SFR, e.g., Twarog 1980; Carlberg et al.
1985: Strobel 1901: see Meusinger ct al.
1985; Strobel 1991; see Meusinger et al.
19901 for a review: Pilvugin Edmunds 1996a.b).. possibly combined with radial inflow across the disc (Carlbere et al.
1991 for a review; Pilyugin Edmunds 1996a,b), possibly combined with radial inflow across the disc (Carlberg et al.
1985. but see Lacey Fall 1985 and Aleusinger ct al.
1985, but see Lacey Fall 1985 and Meusinger et al.
1991). causing metal enrichment to occur in the disc. it has proven no easy task to explain the observed scatter about the relationship in terms of realistic physical processes.
1991), causing metal enrichment to occur in the disc, it has proven no easy task to explain the observed scatter about the relationship in terms of realistic physical processes.
Several mechanisms have been suggested. to be responsible for the observed. spread in metal abundance at any given age (see Carraro et al.
Several mechanisms have been suggested to be responsible for the observed spread in metal abundance at any given age (see Carraro et al.
1998 [or a review): An attempt to study the MIC. outside the solar neighbourhood has been mace by Jonch-Sorensen (1995). who stucied ansihi sample of F and carly C-type stars in six selected. directions within the Galactic disc. and at e-heights twpically below 2 kpc. thereby including both thin and thick clise stars.
1998 for a review): An attempt to study the AMR outside the solar neighbourhood has been made by rensen (1995), who studied an sample of F and early G-type stars in six selected directions within the Galactic disc, and at -heights typically below 2 kpc, thereby including both thin and thick disc stars.
He found a remarkably similar/hin-disc AMI (20.7 kpe) to the local AMI discussed above. with a slope of ~0.04 dex Gyr+ for stars in the age range [rom 210 12 Gyr.
He found a remarkably similar AMR $z < 0.7$ kpc) to the local AMR discussed above, with a slope of $\sim 0.04$ dex $^{-1}$ for stars in the age range from 2 to 12 Gyr.
Alternatively. the AMI for Galactic open clusters has been studied by a number of authors (e... Strobel 1991: Friel Janes 1993: Carraro Chiosi 1994: Carraro et al.
Alternatively, the AMR for Galactic open clusters has been studied by a number of authors (e.g., Strobel 1991; Friel Janes 1993; Carraro Chiosi 1994; Carraro et al.
1998).
1998).
Open clusters are supposed to be good tracers of the Galactic AALR. because their ages. abundances. ancl positions can be estimated to high accuracy. and they represent the old. thin disc population (e.g.. Strobel 1991: Frick Janes 1993).
Open clusters are supposed to be good tracers of the Galactic AMR, because their ages, abundances, and positions can be estimated to high accuracy, and they represent the old thin disc population (e.g., Strobel 1991; Friel Janes 1993).
Although selection cllects related to. among others. the cluster survival mechanisms. their position in the Galactic disc and the possible existence of a spatial metallicity eracdient play a strong role in the determination of the open cluster AAD (e.g. Eriel Janes 1993: Carraro Chiosi 1994: Carraro et al.
Although selection effects related to, among others, the cluster survival mechanisms, their position in the Galactic disc and the possible existence of a spatial metallicity gradient play a strong role in the determination of the open cluster AMR (e.g., Friel Janes 1993; Carraro Chiosi 1994; Carraro et al.
1998). the AALRs for open clusters and stars are essentially the same. including the large scatter about the relationship.
1998), the AMRs for open clusters and stars are essentially the same, including the large scatter about the relationship.
A similar relationship as for the solar neighbourhood stars has been found for the open clusters in the Galactic anticenter direction (Eriel Janes 1993) as well as in the LAIC (o... Olzewski et al.
A similar relationship as for the solar neighbourhood stars has been found for the open clusters in the Galactic anticenter direction (Friel Janes 1993) as well as in the LMC (e.g., Olzewski et al.
1991: Strobel 1991: Eriel Janes 1993). although chemical enrichment has taken place at a slower rate in the LMC than in the Galaxy.
1991; Strobel 1991; Friel Janes 1993), although chemical enrichment has taken place at a slower rate in the LMC than in the Galaxy.
Although it is well-known that the velocity clispersions of disc stars vary with age. the lack of accurate kinematica observations has prevented. the detailed. study of the correlation between the mean stellar age and their vertica velocity. dispersions (referred. to. henceforth).
Although it is well-known that the velocity dispersions of disc stars vary with age, the lack of accurate kinematical observations has prevented the detailed study of the correlation between the mean stellar age and their vertical velocity dispersions (referred to henceforth).
. In addition. a number of discrepant. AVIUss has appeared in the literature (c.e@.. Mayor 1974: Wielen. LOT> Carlhere e al.
In addition, a number of discrepant 's has appeared in the literature (e.g., Mayor 1974; Wielen 1977; Carlberg et al.
1985: Strómmeren 1987: Freeman 1991: Meusinger ct al.
1985; Strömmgren 1987; Freeman 1991; Meusinger et al.
1991: Gommez et al.
1991; Gómmez et al.
1997: see Haywood. Robin Crézzé 1997a.b for a comparison). adding confusion to the already obscure picture.
1997; see Haywood, Robin Crézzé 1997a,b for a comparison), adding confusion to the already obscure picture.
An absorption band at 22.27 qun was first detected on the bright centaur Pholus (Davies. 1993).. and CruiXSijuketal.(1998) present the case that this baud is plausibly due to the presence of methanol. though they »oint οιi that the identification is not unique and that other low molecular weiel ivdrocarbous or photolytic products of methanol might fit the spectrum equally. well.
An absorption band at $\sim$ 2.27 $\mu$ m was first detected on the bright centaur Pholus \citep{1993Icar..102..166D}, and \citet{1998Icar..135..389C} present the case that this band is plausibly due to the presence of methanol, though they point out that the identification is not unique and that other low molecular weight hydrocarbons or photolytic products of methanol might fit the spectrum equally well.
Pholus is oue of the redcdest. οἱyjects iu the solar system. again fiting our picwe of optical colors of irradiated hydrocarbOlls weL
Pholus is one of the reddest objects in the solar system, again fitting our picture of optical colors of irradiated hydrocarbons well.
Based on lower sigua-lo-noise spectra wiil 210yerties which resemble the methanol spectrum of Pholus. he presence of netlalol was sugeestedMOD on the NWBOs 1996 GQ21 al 2002 ΝΕΟΣ (Baruccieal.2008a)..
Based on lower signal-to-noise spectra with properties which resemble the water-ice-plus-methanol spectrum of Pholus, the presence of methanol was suggested on the KBOs 1996 GQ21 and 2002 VE95 \citep{2008ssbn.book..143B}.
To exalnine {1is possibility more closely. we coujne the ]xec& spectra of tese 1 wOOjecs [O Increase slg€)1al-to-oise ail cousider the preseuce of methanol.
To examine this possibility more closely, we combine the Keck spectra of these two objects to increase signal-to-noise and consider the presence of methanol.
While the sigual-Oo-Holse reualus low. the presence of asorption features similar to those ou Pholus is certainly. plausile.
While the signal-to-noise remains low, the presence of absorption features similar to those on Pholus is certainly plausible.
Both objects. like Pholus.« are rect.
Both objects, like Pholus, are red.
A απο[ια of other IXRBOs rave recently been repeted to also have absorption features near 2,27 pm. but. unike Pholus. 1996 C21 aud 2002 VE95. to uot have absorption due to water ice (Baruccietal.20]1)..
A handful of other KBOs have recently been reported to also have absorption features near 2.27 $\mu$ m, but, unlike Pholus, 1996 GQ21 and 2002 VE95, to not have absorption due to water ice \citep{2011Icar..214..297B}.
The signa -1o-nolse|1 tlie spectral regio rare low. so it is difficult to determine if the absorption ealures are ‘eal.
The signal-to-noise in the spectral region are low, so it is difficult to determine if the absorption features are real.
To exauiue the possibility tiab oa 02.27 qnn absorption feature occurs on faint IKBOs. we first examine all of the IxXBCJs and. ceutaurs iu the heck sample (which includes uoue of the yotential 1uethiauol ob.jects from the VLT sauple).
To examine the possibility that a $\sim$ 2.27 $\mu$ m absorption feature occurs on faint KBOs, we first examine all of the KBOs and centaurs in the Keck sample (which includes none of the potential methanol objects from the VLT sample).
We find that a sinall uumber of objects have absorpion features at or near the 2.27 jau methanol absorption line. but no single spectrum is sullicienly reliable in tlis region to assert a positive detection.
We find that a small number of objects have absorption features at or near the 2.27 $\mu$ m methanol absorption line, but no single spectrum is sufficiently reliable in this region to assert a positive detection.
To increase tle we stun the 5»ectra of all of the IKBOs and ceutaurs i1 the Ixeck sample except for those with water ice absorjon at the level of that seen o1 2003 ÀZS1 or deeper and tlose already suspected to contain ujetliauol-like features (1996 C2]| aud 2002 VE95).
To increase the signal-to-noise, we sum the spectra of all of the KBOs and centaurs in the Keck sample except for those with water ice absorption at the level of that seen on 2003 AZ84 or deeper and those already suspected to contain methanol-like features (1996 GQ21 and 2002 VE95).
This combiuec spectrum shows residual absor»t1ous cue to water. as would be i[erred [roii the positive detectlons On nost objects.
This combined spectrum shows residual absorptions due to water, as would be inferred from the positive detections on most objects.
Te only uxijor deviation [rom the water ice s»ectrunm occurs at precisely the wavelength of he feature seen o Pholus aud suspected ou 1996 CQ21 aud 2002 VE95.
The only major deviation from the water ice spectrum occurs at precisely the wavelength of the feature seen on Pholus and suspected on 1996 GQ21 and 2002 VE95.
We conclide that the luethauol-ike feature is iudeed presen at a low-level o1i WBOs even though the featu‘e cannot be reliably identifiedT iu παicual spectra (Figuο 6).
We conclude that the methanol-like feature is indeed present at a low-level on KBOs even though the feature cannot be reliably identified in individual spectra (Figure 6).
No livlesis hias eve ‘been formulated for the sporadic presence of methanol on. IKBOs or cellaurs. o ‘than tc) poit1 out that methiatol is conumon in cometary comae. so IL is exyected to be present iu the interior of Ρο.
No hypothesis has ever been formulated for the sporadic presence of methanol on KBOs or centaurs, other than to point out that methanol is common in cometary comae, so it is expected to be present in the interior of KBOs.
Its yresenee is less expected on the surface of KhBOs. however. as the absorpti features of Ivdrocarbo1 ices ۟ckly degrae under irradiatiou while the remuants turn red (Bruettoetal.2006)..
Its presence is less expected on the surface of KBOs, however, as the absorption features of hydrocarbon ices quickly degrade under irradiation while the remnants turn red \citep{2006ApJ...644..646B}.
Visible metha11ol absorptio1 features suggest that the methanol lias ouly receutly een, exposed at tlie surface. peraps as a result of a collision exposing the subsurface.
Visible methanol absorption features suggest that the methanol has only recently been exposed at the surface, perhaps as a result of a collision exposing the subsurface.
One prediction of this suggestion would be tal the amotnt of inethanol observed would vary as different. faces of the «jbjeet. were observed.
One prediction of this suggestion would be that the amount of methanol observed would vary as different faces of the object were observed.
W‘hile such a test is possible in principle. in practice spectroscopy of these [aint objects is sulficieuv difficult tLat variation would be difficult to prove.
While such a test is possible in principle, in practice spectroscopy of these faint objects is sufficiently difficult that variation would be difficult to prove.
estimate of the difference in velocity between the sides of the bubble would still be zz 20 pc/Myr 2005)..
estimate of the difference in velocity between the sides of the bubble would still be $\approx$ 20 pc/Myr \citep{Maciel_Lago}.
So, we find that at most the galactic shear would stretch this bubble (20Xn)x8Myrs—160(Macielpc, which is much smaller than the 500 pc width of the bubble.
So, we find that at most the galactic shear would stretch this bubble $(20~ {\frac{\rm{pc}}{\rm{Myr}}})\times 8~ \rm{Myrs}= 160~ \rm{pc}$, which is much smaller than the 500 pc width of the bubble.
Therefore, we expect that a bubble would be disturbed by the shear but not torn apart since the upper limit on the shear is still smaller than the bubble.
Therefore, we expect that a bubble would be disturbed by the shear but not torn apart since the upper limit on the shear is still smaller than the bubble.
Even for the simulation without a magnetic field, the center of the bubble only rose 59 pc from the initial explosion height.
Even for the simulation without a magnetic field, the center of the bubble only rose 59 pc from the initial explosion height.
Adding a magnetic field of 44G parallel to the mid-plane pins the cool shell around bubble down and lessens its vertical rise to 23 pc above the explosion height.
Adding a magnetic field of $\mu$ G parallel to the mid-plane pins the cool shell around bubble down and lessens its vertical rise to 23 pc above the explosion height.
However, if the ambient magnetic field is directed perpendicular to the mid-plane then the vertical distribution of the bubble increases.
However, if the ambient magnetic field is directed perpendicular to the mid-plane then the vertical distribution of the bubble increases.
For this case we see hot gas as high as 760 pc above the galactic mid-plane at 8 Myrs, though the explosion went off at 400 pc.
For this case we see hot gas as high as 760 pc above the galactic mid-plane at 8 Myrs, though the explosion went off at 400 pc.